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1607

1607.01526_arXiv.txt

The Fornax dwarf spheroidal galaxy is the most massive satellites of the Milky Way, claimed to be embedded in a huge dark matter halo, and the only among the Milky Way satellites hosting five globular clusters. Interestingly, their estimated masses, ages and positions seem hardly compatible with the presence of a significant dark matter component, as expected in the $\Lambda$ CDM scheme. Indeed, if Fornax would have a CDM halo with a standard density profile, all its globular clusters should have sunk to the galactic centre many Gyr ago due to dynamical friction. Due to this, some authors proposed that the most massive clusters may have formed out of Fornax and later tidally captured. In this paper we investigate the past evolution of the Fornax GC system by using both a recently developed, semi-analytical treatment of dynamical friction and direct $N$-body simulations of the orbital evolution of the globular clusters within Fornax and of Fornax galaxy around the Milky Way. Our results suggest that an ``in-situ'' origin for all the clusters is likely if their observed positions are close to their spatial ones and their orbits are almost circular. Moreover, the Milky Way seems to accelerate the GC decay reducing the decay time of $15\%$. Nevertheless, our results indicate that the GCs survival probability exceeds $50\%$, even in the case of cuspy density profiles. We conclude that more detailed data are required to shed light on the Fornax dark matter content, to distinguish between a cuspy or a cored profile.

The Fornax galaxy is the brightest dwarf spheroidal galaxy (dSph) satellite of the Milky Way (MW) and is, among all of the MW satellites, the only one that hosts five globular clusters, located within $1$ kpc from its center, named Fornax 1,2,3,4 and 5. Moreover, Fornax does not show any bright nucleus. As the other MW satellites, Fornax is believed to reside within a dark matter halo, since its internal velocities are larger than expected from the measurement of the luminous mass \citep{mateo}. The presence of these five, metal poor and old, clusters in Fornax, with ages greater than 10 Gyr \citep{lrs}, represents a still open puzzle because the dynamical friction process (df) as estimated according to the overall galaxy characteristics should have already dragged all the clusters toward the galactic center. Actually, early calculations of the dynamical friction decay time for Fornax lead to few Gyr \citep{tremaine, ohlirich} . This is often referred to as the ``timing problem", which can be posed this way: if the clusters are in the final phase of their orbital decay we are in the unlikely state of looking at them just before their final sink to the galactic center. To study, and hopefully solve, the Fornax puzzle the minimal ingredients are (i) a reliable description of the dynamical friction process in this specific context and (ii) a detailed as possible definition of the phase space profile of the galaxy. With regard to the first point, we remind that the dynamical friction effect is often estimated by mean of the classical Chandrasekhar's formula in its local approximation \citep{Cha43I}, which is well suited to describe the dynamics of massive bodies traveling an extended, isotropic, system, only. It has been indeed proved that Chandrasekhar's approximation fails in more general cases, where more suited approximations have been proposed, as those by \cite{Bin77} and \cite{Pes92} for axisymmetric and triaxial systems, or by \cite{AntMer12} and \cite{ascd14df} for spherical, but cuspy, profiles. On another side, to have a meaningful description of the motion of its clusters, a detailed knowledge of the dynamical structure of Fornax is also required. At this regard, although this galaxy (like all the dSphs) is believed to have a massive dark matter halo (DMH), as expected on the base of the $\Lambda$-CDM paradigm, the density profile compatible with available kinematical data for Fornax does not seem to match the DMH profiles predicted from that paradigm \citep{walker}. Indeed, while CDM predicts the formation of haloes with density profiles scaling as $\rho(r)\propto r^{-1}$ \citep{NFW96}, the kinematic data available for the Fornax clusters suggest flatter density profiles \citep{flores,moore,gilmore,cowsik,jardel}. Many authors provided several explanations for the structure of Fornax and its GCs dynamics. As example, while some authors claim that the galaxy has a flat profile, with a core extended out to $300$ pc from the center \citep{strigari}, others propose that supernovae events could have injected sufficient energy to the environment to remove the DM cusp, leaving a cored profile \citep{pontzen12}. Moreover, due to the recent discovery of shell-like overdensities in that galaxy, some authors have proposed that Fornax is the result of a collision with a smaller galaxy \citep{olsew,coleman04,yozin}. On another side, \cite{angus}, argued that Modified Newtonian Dynamics (MOND, \citep{milgrom83}), can explain the observed dynamical features of the Fornax GCs. \cite{Goerdt} and \cite{Read} showed, via numerical simulations, that in cored profiles a massive body would experience an initial phase in which df is larger than that evaluated with the local approximation formula, going to a second phase in which the test particle stalls. Furthermore, \cite{cole} showed that the timing problem could be solved either by df stall in a cored profile or by the hypothesis that the clusters have formed out of Fornax. The first possibility stands upon the discovery that df becomes unefficient when the infalling object moves on an orbit which encloses a mass nearly equal to the object mass \citep{Gual08,ascd14df}. In this case, the objects reach a nearly stable orbit, never approaching the innermost region of the host galaxy. The second possibility, instead, requires that the GCs have formed in a peripheral region of the MW, and that they have been later tidally captured by Fornax. We stress that the study of the Fornax GC system would help to unveil the mistery hidden in the dynamics of dwarf galaxies, since a convincing explanation of the timing problem would lead to an improvement in the knowledge of the structure of dwarf galaxies, including its dark matter content. In this paper we investigate by mean of high resolution, direct, $N$-body simulations, the effect of the combined tidal field of the host Fornax galaxy and of the MW on the dynamics of the Fornax clusters. By mean of realistic estimates of the df decay times, we could provide a set of possible initial conditions to use in the direct $N$-body simulations. Later on, using reliable models for the Fornax dSph and considering either the case of a cored and a cuspy density profile, we simulated the dynamics of the Fornax GCs, along the motion of Fornax around the Milky Way host. The paper is organized as follows: in Section \ref{model} we present the models for the Galactic tidal field as well as the known properties of the Fornax orbit and the model for the Fornax galaxy and its globular cluster system. Section \ref{df} introduces semi-analytical estimates for the dynamical friction timescales, to place constraints on the cluster initial conditions to use in $N$-body simulations. Section \ref{nb} is devoted to present and discuss results of our direct $N$-body simulations; finally, in Section \ref{end} we draw the conclusions.

\label{end} In this paper we revisited the so called ``timing problem'' for the stellar clusters in the Fornax dwarf spheroidal galaxy. Using two complementary approaches, a simple semi-analytic investigation of the dynamical friction decay times and a more sophisticated and detailed series of numerical simulations for the clusters motion in Fornax as a satellite of the Milky Way, we obtained results which can be summarized as follows: \begin{itemize} \item we found that the missing orbital decay of the Fornax GCs is compatible either with a shallow profile or a steep cusp in the Fornax density profile. This means that a standard CDM density profile cannot be excluded; \item in the extreme hypothesis that the present 3D positions of the GCs coincide with their projected positions, we found that all the clusters should have formed within the Fornax tidal radius, quite independently of the Fornax density profile, even in the case of the most massive cluster (GC3); \item we investigated the gravitational effects induced by the MW tidal field using a series of detailed $N$-body simulations focused on nearly circular orbits. Our results show that in the majority of the investigated cases ($60\%$), the MW tidal field shorten the decay time-scale, leading to its decrease of a factor up to $15\%$; \item on the other hand, we have also found a significant fraction of cases ($35\%$) in which the MW acts against dynamical friction, increasing the decay time, and a small fraction of cases ($5\%$) where the GC is tidally captured by the MW; \item the previous points highlight the importance of both the MW tidal field and of the GCs ICs. Indeed, if the clusters were born in an outer region of Fornax, the MW tidal field tends to slow down their orbital decay but, on the other hand, the MW makes the orbital decay faster for GCs initially moving on orbits within the Fornax scale-radius; \item the MW tidal field induces the formation of tidal tails around Fornax, containing clumps whose surface densities are about 10 times higher than the density of the surrounding tail; \item if the GCs move on nearly circular orbits, there is a wide range of ICs for which they can survive up to a Hubble time, even in the case of a steep Fornax density distribution, thus providing a satisfactory solution to the timing problem and making extremely hard to discern about the shape of possible different mass distributions for Fornax. \end{itemize} In conclusion, this paper shows that the timing problem for the Fornax GCs can be solved even in the case that Fornax has a steep density profile, unless the GCs started moving, at their birth, on nearly radial orbits. Moreover, we have demonstrated that the Galactic gravitational field affects marginally the results, leading in general to shorter decay times in dependence on the IC set. On the other hand, we notice that in few cases, some GCs have been tidally captured by the MW. As a side effect, our results indicate that a standard dark matter mass distribution cannot be completely excluded for Fornax on the base of its GC dynamics.

1607.01526

1607

1607.03009_arXiv.txt

{\salvo{The accretion process in Classical T Tauri Stars (CTTSs) can be studied through the analysis of some UV and X-ray emission lines which trace hot gas flows and act as diagnostics of the post-shock downfalling plasma. In the UV band, where higher spectral resolution is available, these lines are characterized by rather complex profiles whose origin is still not clear.}} {\salvo{We investigate the origin of UV and X-ray emission at impact regions of density structured (fragmented) accretion streams. We study if and how the stream fragmentation and the resulting structure of the post-shock region determine the observed profiles of UV and X-ray emission lines.}} {We model the impact of \salvo{an accretion stream consisting of a series of dense blobs} onto the chromosphere \salvo{of a CTTS through} 2D \salvo{MHD} simulations. We explore different \salvo{levels of stream fragmentation and accretion rates. From the model results, we synthesize C\,IV (1550 \AA) and O\,VIII (18.97 \AA) line profiles.}} {\salvo{The impacts of accreting blobs onto the stellar chromosphere produce reverse shocks propagating through the blobs and shocked upflows. These upflows, in turn, hit and shock the subsequent downfalling fragments. As a result, several plasma components differing for the downfalling velocity, density, and temperature are present altoghether. The profiles of C\,IV doublet are characterized by two main components: one narrow and redshifted to speed $\approx 50$~km s$^{-1}$ and the other broader and consisting of subcomponents with redshift to speed in the range $200 - 400$~km s$^{-1}$. The profiles of O\,VIII lines appear more symmetric than C\,IV and are redshifted to speed $\approx 150$~km s$^{-1}$.}} {\salvo{Our model predicts profiles of C\,IV line remarkably similar to those observed and explains their origin in a natural way as due to stream fragmentation.} }

Classical T Tauri Stars (CTTS) are young stars surrounded by \salvo{an accretion} disk. \salvo{According to the magnetospheric accretion scenario, gas from the disk accretes onto the star through accretion columns \citep{1991Apj...370L..39K}. Several lines of evidence support this picture especially in the optical and infrared bands (e.g. \citealt{1988ApJ...330..350B}).} Also accreting T Tauri Stars show a \salvo{soft X-ray} (0.2-0.8 KeV) excess, with typical lines produced at \salvo{temperatures} $10^5-10^6$ K. \salvo{This excess has been interpreted as due to impacts of the accreting material with the stellar surface where a shock is produced and dissipates the kinetic energy of the downfalling material \citep{1991Apj...370L..39K}. The shock heats the plasma up to temperatures of few million degrees, causing X-ray emission \citep{2002ApJ...567..434K, 2007A&A...465L...5A}. The heated plasma is characterized by density of $n\approx 10^{11}-10^{13}$ cm$^{-3}$ (e.g. \citealt{2007A&A...465L...5A}).} \salvo{The interpretation of the soft X-ray excess in CTTSs in terms of accretion shocks is well supported by hydrodynamic (HD) and magnetohydrodynamic (MHD) models.} {Time-dependent 1D model of radiative accretion shocks in CTTSs provided a first accurate description of the dynamics of the postshock plasma \citep{2008MNRAS.388..357K,2008A&A...491L..17S}. In particular \cite{2008A&A...491L..17S} proposed a one.dimensional hydrodynamic (HD) model of a continuous accretion flow impacting onto the chromosphere of a CTTS, thus assuming the ratio between the thermal pressure and the magnetic pressure $\beta <<1$. Their model reproduced the main features of high spectral resolution X-ray observations of the CTTS MP Mus that was previously interpreted as due to postshock plasma \citep{2007A&A...465L...5A} } \salvo{More recently 2D MHD models of accretion impacts have been studied (\citealt{2010A&A...510A..71O, 2013A&A...557A..69M, 2013A&A...559A.127O}) to explore those cases where the low-$\beta$ approximation cannot be applied (and, therefore, the 1D models cannot be used). These models have shown that the accretion dynamics can be complex with the structure and stability of the impact region of the stream strongly depending on the configuration and strength of the stellar magnetic field. Depending on the magnetic field strength, the atmosphere around the impact region can be also perturbed, leading to accreting plasma leaks at the border of the main stream.} \salvo{Although HD and MHD models of accretion shocks have provided a theoretical support to the hypothesis that the soft X-ray excess in CTTSs originates from impacts of accretion columns onto the stellar surface, several points still remain unclear. Some of these points concern the emission in the UV band arising from impact regions. There is} evidence that \salvo{a significant amount of plasma at $10^{5}$ K (much larger than expected from current models; Costa et al. 2016, submitted)} \salvo{is} produced \salvo{in the accretion process}. \cite{2013ApJS..207....1A} analyzed UV spectra collected with the {\it Hubble Space Telescope} (HST) of 28 T Tauri stars and studied the C\,IV doublet at 1550 \AA. They found that each component of the doublet is described by 2 Gaussian components with different speed and width. \salvo{About half of their sample exhibits line profiles analogous (but with different Doppler shifts) to those of TW Hya that consist of a narrow component redshifted at speeds of $\approx 30$ km s$^{-1}$ (positive speed indicates material that falls into the star) and a broader component centered at $\approx 120$ km s$^{-1}$ and with the redshifted wing reaching $\approx 400$ km s$^{-1}$. The latter component cannot be explained, as post-shock emission, with current models of a continuous accretion stream: assuming a free fall velocity of $\approx 500-600$ km s$^{-1}$ and a strong shock, the velocity of the post-shock plasma cannot be larger than $\approx 100-120$ km s$^{-1}$ \citep{2010A&A...522A..55S}.} \salvo{Recently \cite{2014ApJ...797L...5R} proposed an explanation on the origin of the observed asymmetries and redshifts of UV emission lines in CTTSs. These authors studied the impacts of dense plasma fragments falling back on the surface of the Sun after a violent eruption occurred on 2011 June 7, showing that this phenomenon reproduces on the small scale accretion impacts onto CTTSs (see also \citealt{2013Sci...341..251R}). They modeled the impacts with HD simulations and synthesized the emission in UV and X-ray bands. They found that UV emission may originate from the shocked front shell of the still downfalling fragments, thus producing a broad redshifted component in UV lines up to speeds around $\approx 400$ km s$^{-1}$.} \salvo{In this work we investigate further the scenario of a fragmented stream through MHD modeling. More specifically we study the structure and stability of the post-shock plasma after the impact of a clumpy or fragmented stream onto the stellar surface. We investigate the origin of UV and X-ray emission at impact regions and if and how the stream fragmentation can be responsible of the observed asymmetries and redshifts of UV emission lines in CTTSs. To this end, we developed an MHD model describing an accretion column consisting of several high density blobs which impact onto the chromosphere of a CTTS. We synthesized the C\,IV (1550 \AA) and O\,VIII (18.97 \AA) emission lines, including the effect of Doppler shift due to plasma motion along the line of sight. The paper is structured as follow: in Sect.~\ref{sec2} we describe the model and the synthesis of emission lines; in Sect.~\ref{sec3} we discuss the results of the simulations and the synthesis of emission lines; and finally in Sect.~\ref{sec4} we drawn our conclusions.}

\label{sec4} In this work we investigated the effects of \salvo{stream fragmentation} on C\,IV and O\,VIII emission lines. We developed a model that describes \salvo{a stream composed by a series of dense blobs} impacting onto the surface of a CTTS. Our model \salvo{takes} into account the stellar magnetic field, the gravity, the radiative \salvo{losses} from optically thin plasma and the thermal conduction. \salvo{The aim was to explore if and how the impact of a fragmented stream onto the chromosphere of a CTTS reproduces profiles of C\,IV doublet similar to those observed by \cite{2013ApJS..207....1A}. Our main findings can be summarized as follow:} \begin{itemize} \item \salvo{The impact of a series of dense blobs onto the stellar surface produces a more complex post-shock region than that in the case of a continuous stream. In particular the blob impacts produce strong shocks propagating through the blobs and then upflows after the blobs are fully shocked. The upflows in turn hit and shock the still downfalling blobs, producing a large variety of plasma structures (knots, filaments) differing in density, temperature, and downfalling velocity. These structures are not present in the case of a continuous stream.} \item \salvo{If the stream is fragmented the C\,IV (1550 \AA) lines have a highly asymmetric and broad profiles. The lines split into a narrow and intense component redshifted to speed $\approx 50$~km~s$^{-1}$ and a multitude of faster components redshifted to speeds in the range between 200 and 400~km~s$^{-1}$. The narrow component originates from the post-shock plasma at the base of the accretion column which cools down under the effect of thermal conduction and radiative losses. The fast components originate from thermal instabilities occurring at high altitudes in the shocked stream and from the plasma structures forming during the interaction of upflowing surges and downfalling blobs. This is in agreement with \cite{2014ApJ...797L...5R}.} \item \salvo{The intensity and velocity of the fast components depend on the stream fragmentation: the finer is the fragmentation, the more intense are the fast components. Assuming a more realistic scenario in which few accretion streams with different viewing angles (and therefore different line shifts) are present, the fast components easily merge together to form a single broad component redshifted to speed $\approx 250$~km~s$^{-1}$. A similar result has been found by \cite{2014ApJ...797L...5R} by adopting an HD model to study blob falls on the solar surface. The narrow component at $\approx 50$~km~s$^{-1}$ and the broad component at $\approx 250$~km~s$^{-1}$ are analogous to those found by \cite{2013ApJS..207....1A} for most of the CTTSs of their sample. Thus we interpret the shape of observed C\,IV lines as evidence of density structured or fragmented accretion streams.} \item \salvo{The O\,VIII (18.97 \AA) lines have a symmetric observable profile redshifted to speed ranging between 100 and 200~km~s$^{-1}$. The redshift increases roughly with the level of stream fragmentation: the shift is the smallest in the case of a continuous flow and the largest for a train of small blobs. In any case our model predicts that accretion impacts would produce detectable shifts in the O\,VIII lines.} \item \salvo{As in the case of continuous accretion streams, models of fragmented streams reproduce quite well the luminosity of O\,VIII lines measured in CTTSs \citep{2007A&A...465L...5A,2010ApJ...710.1835B,2012ApJ...752..100A} and, in general, underestimate even by orders of magnitude the luminosity of C\,IV lines. On the other hand, we found that assuming an high level of stream fragmentation is in better agreement with observations. In these models, in fact, many interactions between upflowing surges and downfalling blobs are present which produce a multitude of plasma structures with different density and temperature. Many of them cool down because of radiative losses thus contributing to emission in C\,IV lines. We conclude that the stream fragmentation enhances the emission in the UV band.} \end{itemize} In conclusion, our models reproduce profiles of C IV and O VIII lines remarkably similar to those observed (\citealt{2013ApJS..207....1A}, Argiroffi in preparation) and suggest that the UV emission originates mainly from plasma structures developed as a result of the impact of a density structured or fragmented accretion stream. On the other hand our models predict in general UV luminosities lower than observed. We note that oru models assume that the plasma is optically thin in the whole domain. However, optically thick plasma, as that of the chromosphere and of the unshocked stream, is present around the impact region. This plasma, on one hand, absorb part of the X-ray emission arising from the post-shock plasma \citep{2010A&A...522A..55S,2014ApJ...795L..34B} and, on the other hand, can be heated up to $ \log T (K) \approx 5$ by irradiation by the post-shock plasma, thus contributing to UV emission (Costa et al. 2016, submitted). In this work we did not take into account the absorption by optically thick material and the effect of radiative heating of the unshocked stream by post-shock plasma. In a future work we plan to include these effects on the model to investigate more deeply the origin of UV emission arising from impact regions of fragmented accretion streams.

1607.03009

1607

1607.06240_arXiv.txt

It is known that HLL-type schemes are more dissipative than schemes based on characteristic decompositions. However, HLL-type methods offer greater flexibility to large systems of hyperbolic conservation laws because the eigenstructure of the flux Jacobian is not needed. We demonstrate in the present work that several HLL-type Riemann solvers are provably entropy stable. Further, we provide convex combinations of standard dissipation terms to create hybrid HLL-type methods that have less dissipation while retaining entropy stability. The decrease in dissipation is demonstrated for the ideal MHD equations with a numerical example.

We consider the numerical solution of systems of hyperbolic conservation laws of the form \begin{equation} \label{eq:hypLaw} \pderivative{\vec{q}}{t} + \nabla\cdot\vec{f} = 0, \end{equation} on a domain $\Omega$. For a one-dimensional approximation we divide $\Omega$ into $K$ non-overlapping grid cells $C_i = [x_{i-1/2},\; x_{i+1/2}],\, i= 1,\ldots, K$ which are not necessarily equidistant. In the context of finite volume schemes, hyperbolic equations, such as \eqref{eq:hypLaw}, require a numerical flux function which fully determines the properties of the scheme \cite{LeVeque2002}. The numerical flux function takes as input the left and right value of $\vec{q}$ at the cell interface and solves a local Riemann problem. Smooth initial flows governed by \eqref{eq:hypLaw} may develop discontinuities (\eg shocks) in finite time. Thus, solutions are sought in the weak sense \cite{LeVeque2002}. However, weak solutions are not unique and need to be supplemented with additional admissibility criteria. Following the work of \eg \cite{Winters2016,Tadmor2003}, we use the concept of entropy to construct discretizations that agree with the second law of thermodynamics. That is, the numerical flux function will possess entropy stability, cf. \cite{Tadmor2003} and references therein. In particular, we prove entropy stability for the HLL scheme and present the construction of HLL-type entropy stable numerical flux functions. It is known that HLL-type schemes are more dissipative than upwind schemes. However, HLL-type methods need less information about the eigendecomposition of the flux Jacobian. This is advantageous because the eigenstructure might be computationally expensive or no analytical expression exists, especially for large systems. As such, we consider three standard dissipation terms, namely Lax-Friedrichs (LF), HLL, and Lax-Wendroff (LW) and present two hybrid dissipation terms introduced in \cite{Schmidtmann2016}. We demonstrate that these five schemes are provably entropy stable. % The paper is organized as follows: Sec.~\ref{scn:ESSolvers} provides a brief background on entropy stable numerical fluxes. In Sec.~\ref{scn:Riemann} we show entropy stability for the LF, HLL, and LW dissipation terms. The creation of two new hybrid entropy stable dissipation operators is shown in Sec.~\ref{scn:HybridRiemann}. We demonstrate in Sec.~\ref{sec:NumSim} that the new hybrid numerical flux reduce the overall dissipation in a standard finite volume scheme. Our conclusions and outlook are drawn in the final section.

\label{scn:Conclusion} In this work we constructed two one-parameter families of hybrid entropy stable numerical fluxes. An advantage of the new numerical flux functions is that they remain applicable even when the eigenstructure of the flux Jacobian matrix is unknown. % The derivations and proofs in this work are kept general such that the hybrid entropy stable solvers can be applied to a broad range of non-linear hyperbolic conservation laws. As an example, we applied the novel numerical fluxes to the ideal MHD equations and demonstrated the decreased magnitude of dissipation for the hybrid solvers versus a standard solver. % In the future we plan to apply the hybrid entropy stable Riemann solvers to other complex hyperbolic systems, such as the two-layer shallow water or the regularized 13-Moment Equations of Grad.

1607.06240

1607

1607.07246_arXiv.txt

} Cumulus convecting storms on Saturn are known for their episodic behavior. Convective events on Saturn can be identified visually and through radio noise. The events occur intermittently within a small range of latitudes and are interspersed with quiescent periods that last for a year or longer \citep{Sromovsky_etal_1983, Dyudina_etal_2007, DelGenio_etal_2007}. The largest of these encircle the planet and are sometimes called Great White Spots \citep{Sanchez-Lavega_1994}. These giant storms exhibit a $\sim$30-year periodicity ($\sim$1~Saturn Year), and have been recorded in 1876, 1903, 1933, 1960 and 1990, in which the storms lasted for 26, 150, 44, 41, and 55 days, respectively \citep{Alexander_1962, Sanchez-Lavega_1994}. All documented planet-encircling storms have occurred in the northern hemisphere, alternately erupting in mid-latitude and equatorial regions \citep{Alexander_1962, Sanchez-Lavega_1994}. These storms exhibit some morphological similarities to Jupiter's South Equatorial Belt (SEB) revival events \citep{Sanchez-Lavega_Gomez_1996}; however, their convective areas occupy an area much narrower than the surrounding jets, and move at speeds that reflect the background zonal mean wind. An SEB revival event may last up to 7 months; however, unlike Saturn's planet-encircling storms, its relatively small convective cores do not form a large organized structure. The 1990 equatorial planet-encircling storm of Saturn was an early scientific target of the Hubble Space Telescope \citep{Sanchez-Lavega_etal_1991, Sanchez-Lavega_1994, Beebe_etal_1992, Westphal_etal_1992}. The data collected during the 1990 storm enabled examinations of Saturn's equatorial dynamics. Based on comparing the equatorial wind speeds between Voyager and Cassini measurements, \cite{Perez-Hoyos_Sanchez-Lavega_2006} suggested that the 1990 storm caused a deceleration in the equatorial jet. Numerical simulations by \cite{Sayanagi_Showman_2007} tested the hypothesis that the storm caused a deceleration. They demonstrated that the observed apparent wind change can be explained by the deceleration of the wind caused by the storm in combination with the change in the altitudes of the tracked clouds. This interpretation is consistent with the radiative transfer analysis of \cite{Perez-Hoyos_Sanchez-Lavega_2006}. A new planet-encircling storm started on December~5, 2010. The 2010-2011 outburst was the first planet-encircling storm to be studied from a spacecraft in orbit around Saturn. The Cassini orbiter's Radio and Plasma Wave Science (RPWS) instrument detected radio pulses emitted by lightning discharges in the storm \citep{Fischer_etal_2011Nature}, and the Composite Infrared Spectrometer (CIRS) measured the storm's large-scale heating of the stratosphere and a longitudinal thermal perturbation that is consistent with the formation of a new anticyclonic vortex \citep{Fletcher_etal_2011Science, Fletcher_etal_2012}. Also, ground-based observations captured the development of the large scale disturbance \citep{Sanchez-Lavega_etal_2011Nature}. The mechanism that controls the intermittent appearances and scales of these extreme outbursts is unknown; a better understanding of the conditions that precede these episodes is one step toward unveiling the storms' origin. Here, we analyze Cassini ISS images and RPWS data to investigate the temporal evolution of the 2010-2011 planet-encircling storm. Our observations cover the temporal evolution of the planetary-scale convective disturbance from its beginning on December~5, 2010 until the end of 2011. The rest of our report is organized as follows. Section~2 describes the ISS image sets and the processing applied to them in our study. Section~3 presents the chronological development of the 2010-2011 storm. Section~4 summarizes our findings.

We have reported the evolution of Saturn's latest planet-encircling storm of 2010-2011 using two of the instruments aboard Cassini spacecraft, ISS and RPWS. Our observational coverage enabled us to study the preconditions of the storm, in which we show that the storm erupted out of a previously known feature called the String of Pearls (SoPs). After the storm erupted on December~5, 2010, the outburst grew and engulfed the entire latitude zone. By January 11, 2011, the storm developed a well defined structure with three primary parts as shown in Fig.~5. The western most feature is the particularly bright \emph{head} that propagated at an average rate of $-$2.79$\pm$0.08 \degree~day$^{-1}$ ($-$26.9$\pm$0.8~m~s$^{-1}$), and its longitudinal position is well described by Eq.~\ref{e:head_motion}. A bright body of cloud followed to the east of the head, which is bounded to the east by a new large anticyclonic vortex (AV) that was spawned by the storm. To the east of the AV, a turbulent pattern of clouds formed. The AV maintained a mean drift speed of $-$0.85\degree~day$^{-1}$ ($-$8.4~m$^{~}$s$^{-1}$) between December~24, 2010 and June~14, 2011; during this period, the center of AV did not deviate from Eq.~\ref{e:AV_motion} by more than 10\degree~in longitude (Fig.~17a). We performed cloud-tracking wind measurement on the AV, and showed that its tangential wind speed reached 100~m~s$^{-1}$; we also measured that the relative vorticity within the vortex to be $-$6$\pm$1$\times$10$^{-5}$~s$^{-1}$. With the east-west diameter of 12,000~km, the new anticyclonic vortex became greater in size than any tropospheric vortex previously seen on Saturn. The rapid growth and shrinkage of the new anticyclone is in a stark contrast to the steadiness of anticyclonic vortices on Jupiter. Oval BA today is a result of many dynamical changes over the last $\sim$70 years; however, it has mostly been growing in size through mergers, and the precursor vortices also stayed steady for over 50~years until the mergers. Compared to Jupiter, Saturn has been known to harbor far fewer vortices \citep{Vasavada_etal_2006}. Our analysis suggests that Saturn's inability to maintain even the largest of the vortices contributes to this difference between Jupiter and Saturn. Even though the new Saturnian anticyclone emerged in a highly disturbed region, we do not expect the turbulence to have contributed to its inability to stay large because Jupiter's GRS, for example, resides in the most turbulent latitude of Jupiter and yet it has persisted at least since 1879, and possibly nearly 400 years \citep{Hooke_1665, Beebe_Orton_West_1989}. Lastly, we note that a stratospheric vortex larger in size than the AV emerged in the aftermath of the storm, which persists as of late 2012 \citep{Fletcher_etal_2012DPS}; its future evolution will be of an interest in the context of Saturnian vortex dynamics. The storm underwent a major change in its dynamics after around June~20, 2011 when the storm's head collided with the anticyclone after the anticyclone's longitude trailed 360\degree~longitude behind the core. After the collision, the storm's convective activity displayed a major decline. The SED signals emitted by the storm's lightning activities followed the growth and demise of the cumulus storm observed in the ISS images until the head-AV collision in mid-June. The head-AV collision marked a sharp decline in the SED activities detected by the RPWS. In late August 2011, the SED signal resurged for 9~days, and we identified the source clouds in the ISS images. The RPWS detected two more weak SED activities for the remainder of 2011 during September 30 - October 6, and December 26-28. The storm left the region between 25\degree N and 40\degree N in a highly disturbed state. After the storm, a region that appears particularly dark in CB2 channel emerged between 31\degree N - 38\degree N latitude spanning up to 180\degree~in longitude. After the storm, this region exhibited the lowest albedo on Saturn in the CB2 channel. The region surrounding the dark region exhibited billowing cloud patterns that are similar to mid-latitudes of Jupiter. The storm also altered the zonal mean wind profile of the storm latitudes. Our cloud-tracking measurements reveal that the zonal mean cloud motions accelerated by 35~m$^{~}$s$^{-1}$ to the north of the storm around 38\degree N and decelerated by 30~m$^{~}$s$^{-1}$ to the south around 31\degree N. This change in the zonal wind speed is supported by the tropospheric warming detected by \cite{Achterberg_etal_2012DPS}. The warming is consistent with the latent heat released by the storm. Although the change in the cloud motion could be explained by changes in the cloud altitude and vertical shear, we believe that a real wind speed change was caused by the tropospheric warming. The westward jets at $\pm$35\degree~latitudes have often been the sites of storm activity. Saturn's 35\degree N region harbored several small cumulus storms during the Voyager flybys in 1980-81 \citep{Hunt_etal_1982, Sromovsky_etal_1983}. Earlier in the Cassini mission, similar behaviors, albeit at much smaller scales than the northern-hemisphere planet-encircling storms, were seen in the lightning storms observed in ``Storm Alley,'' at 35\degree S latitude. Like the new storm at 35\degree N, the cumulus events in Storm Alley drifted west and spawned anticyclonic ovals that slowly separated from the source region \citep{Porco_etal_2005, Dyudina_etal_2007}. Unlike in the new storm, the anticyclones drifted west relative to the source, and the source was intermittent. The area of the source region, intensity of the lightning activities, and the size of the anticyclonic vortices were an order of magnitude smaller in the southern storms than in the new northern storm. The five planet-encircling storms previously seen on Saturn alternated their location between equatorial and mid-latitudes \cite{Sanchez-Lavega_1994}. The last event of 1990 was equatorial \citep{Sanchez-Lavega_etal_1991, Sanchez-Lavega_1994, Beebe_etal_1992, Westphal_etal_1992}. The new storm of 2010-2011 followed the previous pattern of events by erupting in mid-latitude. Comparing the earlier storms observed from Earth with the unprecedented details of the present storm revealed by our present study is not straightforward. However, the large-scale structures seen by the ground-based observations of \cite{Sanchez-Lavega_etal_2012gws} are consistent with the results presented here, and give confidence in our ability to compare duration and the spatial structures of the present storm with the past ones. If we take December 5, 2010 as the beginning and June~20, 2011 as the end of the present storm, it lasted for 201~days, longer than any of the five previously recorded storms \citep{Sanchez-Lavega_1994}. Interestingly, the previous longest storm, which lasted for 150 days in 1903, erupted at 30$\pm$2\degree N, a latitude similar to the present storm. On Earth, one of the factors that contribute to the longevity of a thunderstorm is the lack of vertical shear \citep{Emanuel_1994}. The two long-lasting cumulus storms may indicate the lack of vertical shear at their latitudes. Episodic convective outbursts such as the Saturnian planet-circling storms suggest the presence of a mechanism that allows a large build-up of convective available potential energy (CAPE) between the large outbursts. At the depth of 10-bar at the predicted water condensation level, the static stability is believed to be enhanced by the latent heat release \citep{Weidenschilling_Lewis_1973, Sugiyama_etal_2006}. This layer of enhanced static stability could act as a lid for any convective activities that initiate underneath and store CAPE. An episodic cumulus storm can erupt when the internal heat flux deposits sufficient CAPE to break through the water condensation layer; this instability could be triggered spontaneously or modulated by seasonal effects. Perhaps the String of Pearls forms when Rossby waves are excited in the water condensation layer. \cite{Sayanagi_Showman_2007} demonstrated that Rossby waves excited at the 10-bar level can affect the cloud-top level. This possibility can be tested if simultaneous observations by VIMS and ISS reveal vertical offset in the SoPs phase; as those instruments sense different altitudes, comparing their images should reveal the vertical structure of the SoPs. If SoPs is indeed a Rossby wave, the ISS images should show SoPs phase shifted to the west compared to that in the VIMS images. This will be a topic for future investigation. Based on the records of past events, we expect the disturbance to have lasting effects on Saturn's northern hemisphere. As a point of comparison, the 1990 storm left the northern hemisphere of Saturn disturbed for the rest of the decade, and many activities were recorded until 1997 \citep{Sanchez-Lavega_etal_1993a, Sanchez-Lavega_etal_1996, Sanchez-Lavega_etal_1999}. The continuing disturbances of the 2010-2011 storm have been confirmed through ground-based telescopic observations and Cassini RPWS measurements. For example, in early April 2012, the RPWS detected electrostatic discharge signals emitted by a new storm. Cassini CIRS and ground-based infrared observations also continue to show lasting activities in the stratosphere \citep{Hesman_etal_2012DPS, Fletcher_etal_2012DPS, Hesman_etal_2012, Fletcher_etal_2012}, and show a large stratospheric hot spot that has been termed the~``beacon.'' Like the 1990 storm, the aftermath of the latest storm should have an effect that may last for up to a decade, and a continuing monitoring of Saturn from the orbiting vantage point of Cassini spacecraft should reveal further details of the dynamic event. \textbf{Acknowledgement:} Our work was supported by the Cassini-Huygens mission, a cooperative project of NASA, ESA, ASI, managed by JPL, a division of the California Institute of Technology, under a contract with NASA. The authors thank the two anonymous reviewers for their very constructive comments. \label{lastpage} \def\ref@jnl#1{{\jnl@style#1}}% \newcommand\aj{\ref@jnl{AJ}}% \newcommand\araa{\ref@jnl{ARA\&A}}% \newcommand\apj{\ref@jnl{ApJ}}% \newcommand\apjl{\ref@jnl{ApJ}}% \newcommand\apjs{\ref@jnl{ApJS}}% \newcommand\ao{\ref@jnl{Appl.~Opt.}}% \newcommand\apss{\ref@jnl{Ap\&SS}}% \newcommand\aap{\ref@jnl{A\&A}}% \newcommand\aapr{\ref@jnl{A\&A~Rev.}}% \newcommand\aaps{\ref@jnl{A\&AS}}% \newcommand\azh{\ref@jnl{AZh}}% \newcommand\baas{\ref@jnl{BAAS}}% \newcommand\jrasc{\ref@jnl{JRASC}}% \newcommand\memras{\ref@jnl{MmRAS}}% \newcommand\mnras{\ref@jnl{MNRAS}}% \newcommand\pra{\ref@jnl{Phys.~Rev.~A}}% \newcommand\prb{\ref@jnl{Phys.~Rev.~B}}% \newcommand\prc{\ref@jnl{Phys.~Rev.~C}}% \newcommand\prd{\ref@jnl{Phys.~Rev.~D}}% \newcommand\pre{\ref@jnl{Phys.~Rev.~E}}% \newcommand\prl{\ref@jnl{Phys.~Rev.~Lett.}}% \newcommand\pasp{\ref@jnl{PASP}}% \newcommand\pasj{\ref@jnl{PASJ}}% \newcommand\qjras{\ref@jnl{QJRAS}}% \newcommand\skytel{\ref@jnl{S\&T}}% \newcommand\solphys{\ref@jnl{Sol.~Phys.}}% \newcommand\sovast{\ref@jnl{Soviet~Ast.}}% \newcommand\ssr{\ref@jnl{Space~Sci.~Rev.}}% \newcommand\zap{\ref@jnl{ZAp}}% \newcommand\nat{\ref@jnl{Nature}}% \newcommand\iaucirc{\ref@jnl{IAU~Circ.}}% \newcommand\aplett{\ref@jnl{Astrophys.~Lett.}}% \newcommand\apspr{\ref@jnl{Astrophys.~Space~Phys.~Res.}}% \newcommand\bain{\ref@jnl{Bull.~Astron.~Inst.~Netherlands}}% \newcommand\fcp{\ref@jnl{Fund.~Cosmic~Phys.}}% \newcommand\gca{\ref@jnl{Geochim.~Cosmochim.~Acta}}% \newcommand\grl{\ref@jnl{Geophys.~Res.~Lett.}}% \newcommand\jcp{\ref@jnl{J.~Chem.~Phys.}}% \newcommand\jgr{\ref@jnl{J.~Geophys.~Res.}}% \newcommand\jqsrt{\ref@jnl{J.~Quant.~Spec.~Radiat.~Transf.}}% \newcommand\memsai{\ref@jnl{Mem.~Soc.~Astron.~Italiana}}% \newcommand\nphysa{\ref@jnl{Nucl.~Phys.~A}}% \newcommand\physrep{\ref@jnl{Phys.~Rep.}}% \newcommand\physscr{\ref@jnl{Phys.~Scr}}% \newcommand\planss{\ref@jnl{Planet.~Space~Sci.}}% \newcommand\procspie{\ref@jnl{Proc.~SPIE}}% \let\astap=\aap \let\apjlett=\apjl \let\apjsupp=\apjs \let\applopt=\ao

1607.07246

1607

1607.07888_arXiv.txt

We have monitored photometrically the Y0 brown dwarf WISEP J173835.52$+$273258.9 (W1738) at both near- and mid-infrared wavelengths. This $\lesssim 1$ Gyr-old 400~K dwarf is at a distance of 8~pc and has a mass around 5 M$_{\rm Jupiter}$. We observed W1738 using two near-infrared filters at $\lambda \approx 1\,\mu$m, $Y$ and $J$, on Gemini observatory, and two mid-infrared filters at $\lambda \approx 4\,\mu$m, [3.6] and [4.5], on the {\em Spitzer} observatory. Twenty-four hours were spent on the source by {\em Spitzer} on each of June 30 and October 30 2013 UT. Between these observations, around 5 hours were spent on the source by Gemini on each of July 17 and August 23 2013 UT. The mid-infrared light curves show significant evolution between the two observations separated by four months. We find that a double sinusoid can be fit to the [4.5] data, where one sinusoid has a period of $6.0 \pm 0.1$ hours and the other a period of $3.0 \pm 0.1$ hours. The near-infrared observations suggest variability with a $\sim 3.0$ hour period, although only at a $\lesssim 2 \sigma$ confidence level. We interpret our results as showing that the Y dwarf has a $6.0 \pm 0.1$ hour rotation period, with one or more large-scale surface features being the source of variability. The peak-to-peak amplitude of the light curve at [4.5] is 3\%. The amplitude of the near-infrared variability, if real, may be as high as 5 to 30\%. Intriguingly, this size of variability and the wavelength dependence can be reproduced by atmospheric models that include patchy KCl and Na$_2$S clouds and associated small changes in surface temperature. The small number of large features, and the timescale for evolution of the features, is very similar to what is seen in the atmospheres of the solar system gas giants.

There are now more than twenty brown dwarfs known in the solar neighborhood that have effective temperatures ($T_{\rm eff}$) lower than 500~K (Cushing et al. 2011, 2014; Kirkpatrick et al. 2012; Liu et al. 2012; Luhman 2014; Luhman, Burgasser \& Bochanski 2011, Pinfield et al. 2014; Schneider et al. 2015; Tinney et al. 2012). These have been classified as Y dwarfs (Cushing et al. 2011, Kirkpatrick et al. 2012). Evolutionary models show that for $300 \leq T_{\rm eff}$~K $\leq 500$ and 0.2 $\leq$ age~Gyr $\leq$ 10 (as appropriate for the solar neighborhood) the mass range is 2 -- 30 Jupiter masses (Saumon \& Marley 2008). Hence this population of isolated brown dwarfs has a mass that is very planet-like. Our group has an ongoing program measuring the photometric variability of Y dwarfs. For warmer brown dwarfs variability is usually associated with inhomogeneous or variable cloud structure in the atmosphere (e.g. Radigan et al. 2012). For Y0 and Y1 dwarfs with $T_{\rm eff} \approx 400$~K the atmospheres are generally cloud-free, because most of the atmosphere is too cold for chloride or sulphide clouds, and too warm for water or ammonia clouds (Burrows et al. 2003; Morley et al. 2012, 2014); in fact cloud-free models can reproduce Y dwarf observations (Leggett et al. 2015, 2016). Nevertheless variability may be seen at wavelengths where flux is emitted from very high and cold or low and warm layers where condensates can be present (Morley et al. 2012), or variability may be seen due to temperature variations across the brown dwarf surface (Showman \& Kaspi 2013). In our first paper (Cushing et al. 2016, hereafter Cu16) we show that the Y0.5(pec) brown dwarf WISEPC J140518.40$+$553421.5 (W1405, Cushing et al. 2011) is variable at mid-infrared wavelengths. W1405 was observed with {\em Spitzer} using the IRAC camera (Fazio et al. 2004) in the [3.6] and [4.5] filters. Variability was evident at [4.5] in the first epoch and at both [3.6] and [4.5] in the second epoch. The second-epoch light curves have a period of about 8.5 hr, and semi-amplitudes of 3.5\%. In the current paper we present the detection of variability at mid-infrared wavelengths in another Y0, WISEP J173835.52$+$273258.9 (W1738, Cushing et al. 2011). We also present the tentative detection of variability at near-infrared wavelengths, at the $\lesssim 2 \sigma$ confidence level. We extend to lower limits the work of Rajan et al. (2015) who excluded any $J$-band variability larger than 20\% for this brown dwarf. Leggett et al. (2016) compares near-infrared spectra and photometry for W1738 to recent models which include chemical disequilibrium driven by vertical transport (Tremblin et al. 2015). It was necessary to include mixing in order to reproduce the observations, and a cloud-free solar metallicity model with $T_{\rm eff} = 425 \pm 25$K and log $g = 4.0 \pm 0.25$ fit the data best. This temperature and relatively low gravity imply that W1738 is a 3 -- 9 Jupiter mass object with an age of 0.15 -- 1 Gyr. Table 1 lists properties of W1738. In \S 2 we present new observations of W1738 obtained with {\em Spitzer} and IRAC, and Gemini Observatory and its near-infrared imager NIRI (Hodapp et al. 2003). We obtained two epochs of [3.6] and [4.5] data, separated by four months, as well as two epochs of near-infrared $Y$ and $J$ data, obtained between the {\em Spitzer} observations and separated by one month. \S 3 presents our analysis of the data, which we discuss in \S 4. Our conclusions are given in \S 5.

We obtained 12 hours of continuous {\em Spitzer} data on the Y0 brown dwarf W1738 at [3.6], followed by another 12 hours at [4.5]. Two sets of data were obtained four months apart, on June 30 and October 30 2013. We also obtained interspersed Gemini $Y$ and $J$ data on W1738, with about 1.4 hours on-source at $Y$ and 2.3 hours on-source at $J$, on two occasions separated by about one month, July 17 and August 23 2013. Fourier and Lomb-Scargle analyses of the mid-infrared [4.5] data suggested the presence of three and six hour periods in the data. We use a probabilistic method to fit double sinusoids to the [4.5] data (the near-infrared data covers a short time span and the [3.6] data are too noisy). We constrain the second sinusoid to have half the period of the first, but allow amplitude and phase to vary. We find sinusoids with periods of 5.8 $\pm$ 0.1 hours and 6.13 $\pm$ 0.08 hours, and half those values, reproduce the observed [4.5] light curves on the two epochs well. The amplitudes range from 0.3\% to 1.1\%, leading to peak-to-peak variability of 3\%. The shorter time-span near-infrared data was inspected visually only. The data suggest that W1738 is also variable in the near-infrared, although only at the $\lesssim 2 \sigma$ confidence level. If real, the implication is that this Y0 dwarf varied by 10 -- 30\% in $Y$ and 5 -- 15 \% in $J$, peak-to-peak, with a period of about three hours, at two epochs during 2013. The observations are consistent with W1738, a 5 Jupiter-mass $T_{\rm eff} \approx 400$ K Y-type brown dwarf, being seen nearly-equator on, having a rotation period of $6.0 \pm 0.1$ hours, and having one or more large surface features which give rise to the variability. The features evolve over timescales of months. The observed variability at $\lambda \sim 4\,\mu$m is likely due to thermal variations caused by atmospheric circulation, while the larger variations at $\lambda \sim 1\,\mu$m, if real, may be due to the presence of patchy clouds of KCl and Na$_2$S in the lower regions of the atmosphere.

1607.07888

1607

1607.08820_arXiv.txt

Our understanding of the universe relies mostly on electromagnetism. As photons are the messengers, fundamental physics is concerned in testing their properties. Photon mass upper limits have been earlier set through pulsar observations, but new investigations are offered by the excess of dispersion measure (DM) sometimes observed with pulsar and magnetar data at low frequencies, or with the fast radio bursts (FRBs), of yet unknown origin. Arguments for the excess of DM do not reach a consensus, but are not mutually exclusive. Thus, we remind that for massive electromagnetism, dispersion goes as the inverse of the frequency squared. Thereby, new avenues are offered also by the recently operating ground observatories in 10-80 MHz domain and by the proposed Orbiting Low Frequency Antennas for Radio astronomy (OLFAR ). The latter acts as a large aperture dish by employing a swarm of nano-satellites observing the sky for the first time in the 0.1 - 15 MHz spectrum. The swarm must be deployed sufficiently away from the ionosphere to avoid distortions especially during the solar maxima, terrestrial interference and offer stable conditions for calibration during observations.

For many years, many of us have worked, see {\it e.g.}, \cite{spalliccietal2005}, for the opening of one of the gravitational wave windows on the universe, that has finally occurred \cite{abbottetal2016}, while for other windows we still have to wait \cite{spallicci2013}. Nevertheless, we can genuinely state that even when gravitational wave information will be exploited by ground or space laser interferometers, or by the pulsar timing array, our understanding of the large scale universe will be largely (exception made for neutrinos and cosmic rays) based on electromagnetic observations of the four interactions that rule the physical world. As photons are the messengers, fundamental physics has the concern of testing the foundations of electromagnetism, while astrophysics the task of interpreting the universe accordingly. { Furthermore, while alternative theories to general relativity, including those based on massive gravitons, are also conceived for solving the questions raised by the dark universe or to couple gravity with the other interactions, less effort is deployed for studying alternative electromagnetism. But electromagnetism at large scales may differ from the Maxwellian conception of the nineteenth century and thereby contribute to solve some of the riddles in contemporary physics and cosmology}. Since neutrinos have been declared massive, while waiting for the elusive graviton, and leaving aside gluons that are not observed as free particles, the photon is the sole particle of the standard model to be massless. Finally, massive photons would manifest themselves through delays of low frequency electromagnetic signals, being the sub-MHz part of the spectrum yet unexplored \cite{lacki2010}. Thus, this work addresses the potential of the low frequency region to test the foundations of physics. We shall use SI units throughout the paper.

We have focused on the interest raised by the newly operating and future ground and space detectors at very low radio frequencies. The foundations of electromagnetism could be tested, confirming our beliefs or else contributing to establish new physics. Studies on dispersion appear crucial to unveil the causes of the excess of DM, and disentangle the measure of DM from distance. Low frequency observatories placed faraway from the Earth will be an essential aid to such studies, and will possibly set competitive limits to the photon mass. Meanwhile, theoretical investigations on the plausibility and implications of massive photons is to be pursued. In the context of Standard Model Extensions (SMEs), four general classes of Super Symmetry (SuSy) and Lorentz Symmetry (LoSy) breaking were analysed, leading to observable imprints at our energy scales \cite{bodshnsp2016a}. The photon dispersion relations show a non-Maxwellian behaviour for the CPT (Charge-Parity-Time reversal symmetry) odd and even sectors. The group velocities exhibit also a directional dependence with respect to the breaking background vector (odd CPT) or tensor (even CPT). In the former sector, the group velocity may decay following an inverse squared frequency behaviour. Thus, a massive and gauge invariant Carroll-Field-Jackiw photon term in the Lagrangian has been extracted and the induced mass shown to be proportional to the breaking vector. The latter is estimated by ground measurements and leads to a photon mass upper limit of $10^{-19}$ eV or $2 \times 10^{-55}$ kg. Implications for cosmology by non-linear and massive photon theories, and generally non-Maxwellian behaviour have not yet been evaluated adequately. Finally, laboratory experiments should be pursued, investigating in all directions, including the search of frequency shifts, incidentally using the same equipment to set upper limit to the Hubble parameter at small scale \cite{shamirfox1967,bonnor1999,dumin2012,priceromano2012,kopeikin2015}.

1607.08820

1607

1607.06710_arXiv.txt

{The atmospheres of extrasolar planets are thought to be built largely through accretion of pebbles and planetesimals. Such pebbles are also the building blocks of comets. The chemical composition of their volatiles are usually taken to be inherited from the ices in the collapsing cloud. However, chemistry in the protoplanetary disk midplane can modify the composition of ices and gases.} {To investigate if and how chemical evolution affects the abundances and distributions of key volatile species in the midplane of a protoplanetary disk in the 0.2--30 AU range. } {A disk model used in planet population synthesis models is adopted, providing temperature, density and ionisation rate at different radial distances in the disk midplane. A full chemical network including gas-phase, gas-grain interactions and grain-surface chemistry is used to evolve chemistry in time, for 1 Myr. Both molecular (inheritance from the parent cloud) and atomic (chemical reset) initial conditions are investigated.} {Great diversity is observed in the relative abundance ratios of the main considered species: \ce{H2O}, CO, \ce{CO2}, \ce{CH4}, \ce{O2}, \ce{NH3} and \ce{N2}. The choice of ionisation level, the choice of initial abundances, as well as the extent of chemical reaction types included are all factors that affect the chemical evolution. The only exception is the inheritance scenario with a low ionisation level, which results in negligible changes compared with the initial abundances, regardless of whether grain-surface chemistry is included. The grain temperature plays an important role, especially in the critical 20-28 K region where atomic H no longer sticks long enough to the surface to react, but atomic O does. Above 28 K, efficient grain-surface production of \ce{CO2} ice is seen, as well as \ce{O2} gas and ice under certain conditions, at the expense of H$_2$O and CO. \ce{H2O} ice is produced on grain surfaces only below 28~K. For high ionisation levels at intermediate disk radii, \ce{CH4} gas is destroyed and converted into CO and \ce{CO2} (in contrast with previous models), and similarly \ce{NH3} gas is converted into N$_2$. At large radii around 30 AU, \ce{CH4} ice is enhanced leading to a low gaseous CO abundance. As a result, the overall C/O ratios for gas and ice change significantly with radius and with model assumptions. For high ionisation levels, chemical processing becomes significant after a few times $10^{5}$~yrs.} {Chemistry in the disk midplane needs to be considered in the determination of the volatile composition of planetesimals. In the inner $<$30 AU disk, interstellar ice abundances are preserved only if the ionisation level is low, or if these species are included in larger bodies within $10^{5}$~yrs.}

\label{intr} The discovery of more than 2000 extrasolar planets by the radial velocity and transiting techniques \citep[e.g.,][]{udry2007,borucki2011,batalha2013,fischer2014} has signaled the next phase in exoplanet research: the characterization of their atmospheres. Simple molecules such as CO, \ce{H2O} and perhaps \ce{CO2} and \ce{CH4} are being detected in a growing number of exoplanet atmospheres \citep[e.g.,][]{seager2010,snellen2010,birkby2013,fraine2014,crossfield2015,sing2016}. These atmospheres are thought to be built up largely by the accretion of pebbles and planetesimals in the natal protoplanetary disk (see \citealt{johansen2014} and \citealt{benz2014} for reviews), hence the atmospheres should reflect the chemical composition of the disk. There are two different views on how to treat the chemistry in the midplanes of disks, depending on the scientific focus and heritage. Planet formation and population synthesis models \citep[e.g.,][]{ida2004,ida2008,alibert13} consider multiple physical effects taking place in a protoplanetary disk, such as gravitational interactions between bodies, orbital excitation and eccentricity damping, gas drag, accretion of material onto planets, and planet migration in the gaseous disk. Hence, there is a high degree of physical complexity and detail to planet formation processes in these simulations. However, these models do not contain any detailed chemistry. Either they simply use the observed chemical abundances in interstellar ices and assume that these abundances are preserved during disk evolution, or they assume that thermodynamic equilibrium is attained so that chemical abundances are controlled by temperature and pressure only \citep[e.g.,][]{mousis2010,johnson2012, moses2013,marboeuf14,thiabaud2015gascomp}. The main observational test is through statistical comparisons with the observed populations of exoplanets and their predicted compositions. The alternative view starts from detailed physico-chemical models of protoplanetary disks which are closely linked to, and tested by, a wide variety of astronomical observations (see reviews by \citealt{williams11} and \citealt{armitage2011}). Starting from an assumed (static) surface density distribution, scale height and disk flaring, such models first determine the temperature structure of dust and gas heated by the central star through calculation of the full radiative transfer of the dust and the thermal balance of the gas \citep[e.g.,][]{dullemond2007,nomura2005,woitke2009,bruderer2013}. This physical model is then coupled with an extensive gas-grain chemistry network to solve the kinetic chemistry equations at each point in the disk and compute the chemical composition of the gas and ice as a function of time \citep[e.g.,][]{bergin2007,furuya2014,cleeves2014water,reboussin2015,walsh2012,walsh15}. Since planets are formed in the midplanes of disks, it is particularly important to consider the composition and evolution in the midplanes. To what extent is the initial chemical composition of material that is accreted onto a protoplanetary disk preserved, and what happens to the material after it reaches the midplane of the disk, i.e., to what extent is it reset \citep{visser2009,pontoppidan2014}? Does planetesimal formation happen so fast that ices are incorporated into large bodies early on in the evolution, preventing further chemical processing \citep{zhang2015}? Additional clues to the chemical evolution in disks come from the observations of comets in our own solar system \citep{charnley11}. Cometary records suggest that the chemical composition of the pre-solar nebula has been at least partially preserved in the comet-forming zone throughout its lifetime, pointing to little or no chemical processing. However, the original composition of the material that was present in the protoplanetary disk around the Sun when it formed remains unknown, and studies of other disks are needed to provide a framework for our own solar system. Particularly interesting are the recent results from the ESA {\it Rosetta} mission finding significant amounts of \ce{O2} in comet 67P/C-G \citep[see][]{bieler15}, with similarly high O$_2$ abundances inferred for comet Halley from a re-analysis of the {\it Giotto} data \citep{rubin2015}. Abundances as high as a few~\% of solid \ce{O2} with respect to solid \ce{H2O} are not yet fully understood. Lastly, the deuteration of water and organics also provides insight into the history of the pristine material from the ISM \citep[see][]{ceccarelli2014}. \begin{figure*} \subfigure[][]{\includegraphics[width=0.5\textwidth]{T_N_profile_corrected.pdf}\label{temp_dens}} \subfigure[][]{\includegraphics[width=0.5\textwidth]{cr_level_16.pdf}\label{cr_level}} \caption{(a) The temperature $T(R)$ in K (blue) and number density $n(R)$ in cm$^{-3}$ (red) profiles for the disk midplane. The solid blue line indicates the adjusted temperature profile (as described in the text). The original temperature profile from \citet{alibert13} beyond 5.2~AU is indicated by the dashed blue line. (b) The ionisation rate $\zeta(R)$ in s$^{-1}$ adopted for the disk midplane. The red dashed line depicts the contribution to the ionisation rate from short-lived radionuclides (SLRs) only. The blue dotted line is the contribution from external cosmic rays (CRs) only. The solid green line represents the total ionisation rate (SLRs and CRs) as a function of radial distance.} \label{simple_c} \end{figure*} In their planet population synthesis models, \citet{marboeuf14} assumed the initial chemical abundances to be inherited directly from the interstellar ices observed in dense interstellar clouds. A set of eight volatile molecules (\ce{H2O}, CO, \ce{CO2}, \ce{CH4}, \ce{H2S}, \ce{CH3OH}, \ce{N2}, and \ce{NH3}, species also considered in this work) were homogeneously distributed in their model disk midplane with relative ratios consistent with interstellar ice observations \citep{gibb2004,oberg2011ices,boogert2015}. Depending on the physical conditions in different parts of the midplane, as well as the sublimation temperatures of the species, these molecules could then either be assigned to the gas or ice, with the threshold set by the icelines of the species. Icelines (or snowlines) mark the radius in the disk midplane beyond which species exist solely in ice form and are thus depleted from the gas. This occurs at the radius where the accretion rate onto grain surfaces (or freezeout) exceeds the desorption rate from grain surfaces due to the negative temperature gradient in the midplane. The relative rates of these processes are very strong functions of temperature leading to a narrow transition region from gas to ice (moving outwards in radius). The position of the midplane iceline for a particular species will depend on its volatility (i.e., its binding energy). \citet{marboeuf14} do not consider any chemical reactions in their models, besides freezeout and desorption. The positions of the icelines are important because they determine which species are gas and ice at any location in the disk, and thus which material is available to build larger bodies (solids only). If, for example, a giant planet is forming in the disk, the composition of its core will reflect the ice compositions at the different positions in the disk through which the forming planet has moved. The composition of the planet's atmosphere, on the other hand, will reflect the gas composition at the position where the planet becomes massive enough to accrete an atmosphere onto its surface from the surrounding gas in the disk. Moreover, accretion of icy bodies may still pollute the atmosphere. These pebbles and planetesimals migrate through the disk due to radial drift and may therefore have originated at larger radii. Depending on the pebble and planetesimal sizes, the migration of these objects also affects the location of the icelines \citep[see, e.g.,][]{piso2015}. Particularly important is the C/O ratio of the solid and gaseous material in the disk \citep{oberg2011co}. The ratio depends not only on the different volatilities of the chemical species but also on their production or destruction as a consequence of chemical processing. Since \ce{H2O} and \ce{CO2} (which are both O rich) freeze out at higher temperatures than species that are more C rich, such as \ce{CH4} and CO, the C/O ratio depends on both the physical structure and chemistry in the disk. Ultimately, the chemical composition of a planet's core and its atmosphere may thus differ depending on the history of the disk, the formation location of the planet, and any subsequent migration. To address these questions, we use a physical disk model, in particular its midplane temperature and density, which is the same as that considered in the \citet{marboeuf14} population synthesis models. We compute the abundances of chemical species with time using a comprehensive chemical network and different sets of assumptions (see below) to investigate the degree to which chemical evolution/processing affects the resulting abundances of key volatiles in the disk midplane. The sensitivity of our results to the choice of (i) initial chemical abundances (parent cloud inheritance or chemical reset), (ii) the physical conditions (in particular ionisation level), and (iii) the types of chemical reactions included in the model, are also investigated, with details provided in Sect.~\ref{methods}. This generates eight different simulations, the results of which are presented in Sect.~\ref{results}. Sect.~\ref{discussion} discusses the validity of the inheritance and reset scenarios, the implications for planet formation, and the extent to which the results hold for other disk models. Sect.~\ref{conclusions} summarises the conclusions from this work. \begin{table} \caption{Initial abundances (with respect to H$_{\rm nuc}$) for atomic and molecular initial abundances setups. The binding energies $E_{b}$ for all species are also listed.} % \centering % \begin{tabular}{l l c r} % \hline\hline % Species & Atomic & Molecular &$E_{b}$[K]\\ % \hline % \\ H & 9.1$\times 10^{-5}$ & 5.0$\times 10^{-5}$ & 600\\ % He & 9.8$\times 10^{-2}$ & 9.8$\times 10^{-2}$ & 100\\ \ce{H2} & 5.0$\times 10^{-1}$ & 5.0$\times 10^{-1}$ & 430\\ N & 6.2$\times 10^{-5}$ & & 800\\ O & 5.2$\times 10^{-4}$ & & 800\\ C & 1.8$\times 10^{-4}$ & & 800\\ S & 6.0$\times 10^{-6}$ & & 1100\\ \ce{H2O} & &3.0$\times 10^{-4}$ & 5770\\ CO & &6.0$\times 10^{-5}$ & 855\\ \ce{CO2} & &6.0$\times 10^{-5}$ & 2990\\ \ce{CH4} & &1.8$\times 10^{-5}$ & 1090\\ \ce{N2} & &2.1$\times 10^{-5}$ & 790\\ \ce{NH3} & &2.1$\times 10^{-5}$ & 3130\\ \ce{CH3OH} & &4.5$\times 10^{-5}$ & 4930\\ \ce{H2S} & &6.0$\times 10^{-6}$ & 2743\\ \ce{O2} & &0 & 1000\\ \ce{HCN} & &0 & 3610\\ \ce{NO} & &0 & 1600 \\ \hline % \label{init_abun} \end{tabular} \end{table}

\label{conclusions} The models presented in this work have examined the importance of kinetic chemistry on the molecular composition (gas and ice) in protoplanetary disk midplanes. The main conclusions are listed below. \begin{itemize} \item The disk midplane composition reflects that of interstellar ices only for the case of low ionisation (SLRs only) in the inheritance scenario. The partitioning between gas and ice is determined solely by iceline positions, as is assumed in the planet population synthesis models. The inclusion of grain-surface chemistry has a negligible effect. \item Assuming a higher rate of ionisation (SLRs plus CRs) and inheritance leads to an increase in the abundance of \ce{CH4} ice beyond its iceline, and a significant depletion of gas-phase \ce{CH4} in the critical region between 1 and 15~AU. Cosmic-ray-induced chemistry enables the release of free carbon from CO in the outer disk ($>15$~AU) which is incorporated into gas-phase methane which freezes out. This naturally leads to low gas-phase CO abundances as is observed in some disks. On the other hand, cosmic-ray-induced chemistry efficiently destroys methane gas in the \ce{CH4}-poor region. The conclusion holds whether grain-surface chemistry is included or not. \item When grain-surface chemistry is considered, the \ce{CO2} ice to \ce{H2O} ice ratio is increased, with \ce{CO} and \ce{CH4} gas destroyed at the expense of an increase in \ce{CO2} ice. The critical reactions are the photodissociation of \ce{H2O} ice to form \ce{OH} radicals within the ice mantle, which subsequently react with CO to form \ce{CO2} ice. This reaction is able to proceed faster than H~+~OH recombination within $\approx 20$~AU because the very volatile H atoms are quickly lost to the gas phase for grains warmer than 20~K. Beyond this radius, the reformation of water ice wins. The partitioning between \ce{N2} and \ce{NH3} is similarly affected. \item For the extreme reset scenario in which all elements are initially in atomic form, the picture changes significantly. Without grain-surface chemistry, gas-phase CO, \ce{O2}, and atomic oxygen are the main carbon and oxygen-bearing species beyond 1~AU. The chemistry does not have sufficient time to incorporate all available initial elemental oxygen into molecules by $10^6$~yrs. Gas-phase water and \ce{CO2} do form and subsequently freeze out, albeit achieving much lower abundances than in the inheritance scenario. A higher ionisation level helps to increase the production of \ce{H2O} and \ce{CH4}. \item With grain-surface chemistry, the abundances of \ce{H2O} and \ce{CO2} ice increase significantly, demonstrating the absolute necessity of grain-surface chemistry for the synthesis of these two dominant ice components. The final abundance ratios reached in the very outer disk (30~AU) for \ce{H2O} and \ce{CO2} are similar to those for the inheritance scenario, regardless of the ionisation level; however, the higher ionisation level does impede the abundance of \ce{O2} ice at the levels seen in comets and enables a conversion from CO to \ce{CH4}. \item In the reset scenario, species other than the main considered volatiles are also produced in non-negligible quantities: \ce{O2}, \ce{HCN}, and \ce{NO}. The higher ionisation level generally helps the production of HCN and NO (at the expense of \ce{N2} and \ce{NH3}) and impedes the survival of \ce{O2} ice. \item The inclusion of chemistry has a significant impact on the C/O ratio of both gas and ice in the planet-forming region which is expected to influence the resulting composition of forming planet(esimal)s. The ices remain, on the whole, dominated by oxygen (i.e., C/O $< 1$). For both inheritance cases, the gas is carbon rich relative to the canonical value; however only for the low ionisation case is there a reservoir of gas-phase material with C/O $>1$. For both reset scenarios, the ice becomes more carbon rich than the gas, which is opposite to the inheritance case. \end{itemize} The results presented here show that under certain conditions, highlighted above, chemistry can have a profound effect on the composition of the planet-forming material in disk midplanes. Chemistry influences the partitioning of elemental carbon, oxygen, and nitrogen, into molecules of differing volatilities, such that the positions of ice lines alone, are not necessarily adequate for determining the ratio of C/O in neither the gas, nor the ice. Only under the extreme case of full inheritance and low ionisation, are the elemental ratios determined solely by the positions of icelines. This conclusion is also similar to that for the assumption that ices are already locked up in larger bodies by $\approx 10^{5}$~yr. The work presented here follows the time evolution of chemistry in a static protoplanetary disk which is the simplest physical case. In reality, disk conditions evolve with time, at the same time as planetesimals are forming and migrating within the midplane. Future plans include determining the influence of chemistry in an evolving protoplanetary disk (where the density, temperature, and ionisation rate also vary with time), and to couple the outputs of these models with planet formation tracks to determine, in a quantitative manner, the influence on the resulting composition of gas-giant planetary atmospheres.

1607.06710

1607

1607.08619.txt

Mapping Nearby Galaxies at Apache Point Observatory (MaNGA) is an optical fiber-bundle integral-field unit (IFU) spectroscopic survey that is one of three core programs in the fourth-generation Sloan Digital Sky Survey (SDSS-IV). With a spectral coverage of 3622 -- 10,354 \AA\ and an average footprint of $\sim 500$ arcsec$^2$ per IFU the scientific data products derived from MaNGA will permit exploration of the internal structure of a statistically large sample of 10,000 low redshift galaxies in unprecedented detail. %, and also provide %a valuable legacy data set in the era of next-generation high redshift surveys. Comprising 174 individually pluggable science and calibration IFUs with a near-constant data stream, MaNGA is expected to obtain $\sim$ 100 million raw-frame spectra and $\sim 10$ million reduced galaxy spectra over the six-year lifetime of the survey. In this contribution, we describe the MaNGA Data Reduction Pipeline (DRP) algorithms and centralized metadata framework that produces sky-subtracted, spectrophotometrically calibrated spectra and rectified 3-D data cubes that combine individual dithered observations. For the 1390 galaxy data cubes released in Summer 2016 as part of SDSS-IV Data Release 13 (DR13), we demonstrate that the MaNGA data have nearly Poisson-limited sky subtraction shortward of $\sim$ 8500 \AA\ and reach a typical $10\sigma$ limiting continuum surface brightness $\mu = 23.5$ AB arcsec$^{-2}$ in a five arcsec diameter aperture in the $g$ band. The wavelength calibration of the MaNGA data is accurate to $5$ \kms\ rms, with a median spatial resolution of 2.54 arcsec FWHM (1.8 kpc at the median redshift of 0.037) and a median spectral resolution of $\sigma = 72$ \kms.

Over the last twenty years, multiplexed spectroscopic surveys have been valuable tools for bringing the power of statistics to bear on the study of galaxy formation. Using large samples of tens to hundreds of thousands of galaxies with optical spectroscopy from the Sloan Digital Sky Survey \citep{york00,dr1.ref} for instance, studies have outlined fundamental relations between stellar mass, metallicity, element abundance ratios, and star formation history \citep[e.g.,][]{kauffmann03,tremonti04,thomas10}. However, this statistical power has historically come at the cost of treating galaxies as point sources, with only a small and biased region subtended by a given optical fiber contributing to the recorded spectrum. % Realizing limitations of such work though; move to IFU spectroscopy. Other major IFU works. As technology has advanced, techniques have been developed for {\it imaging spectroscopy} that allow simultaneous spatial and spectral coverage, with correspondingly greater information density for each individual galaxy. Building on early work by (e.g.) \citet{colina99} and \citet{dezeeuw02}, such integral-field spectroscopy has provided a wealth of information. In the nearby universe for instance, observations from the DiskMass survey \citep{bershady10} have indicated that late-type galaxies tend to have sub-maximal disks \citep{bershady11}, while Atlas-3D observations \citep{cappellari11a} showed that early-type galaxies frequently have rapidly-rotating components \citep[especially in low density environments;][]{cappellari11b}. In the more distant universe, integral-field spectroscopic observations have been crucial in establishing the prevalence of high gas-phase velocity dispersions \citep[e.g.,][]{fs09,law09,law12,wisnioski15}, giant kiloparsec-sized ÒclumpsÓ of young stars \citep[e.g.,][]{fs11}, and powerful nuclear outflows \citep{fs14} that may indicate fundamental differences in gas accretion mechanisms in the young universe \citep[e.g.,][]{dekel09}. More recently, surveys such as CALIFA \citep[Calar Alto Legacy Integral Field Area Survey,][]{sanchez12,califadr2}, SAMI \citep[Sydney-AAO Multi-object IFS,][]{croom12,allen15}, and MaNGA \citep[Mapping Nearby Galaxies at APO,][]{bundy15} have begun to combine the information density of integral field spectroscopy with the statistical power of large multiplexed samples. As a part of the 4th generation of the Sloan Digital Sky Survey (SDSS-IV), the MaNGA project bundles single fibers from the Baryon Oscillation Spectroscopic Survey (BOSS) spectrograph \citep{smee13} into integral-field units (IFUs); over the six-year lifetime of the survey (2014-2020) MaNGA will obtain spatially resolved optical+NIR spectroscopy of 10,000 galaxies at redshifts $z \sim 0.02 - 0.1$. In addition to providing insight into the resolved structure of stellar populations, galactic winds, and dynamical evolution in the local universe \citep[e.g.,][]{belfiore15,li15,wilkinson15}, the MaNGA data set will be an invaluable legacy product with which to help understand galaxies in the distant universe. As next-generation facilities come online in the final years of the MaNGA survey, IFU spectrographs such as TMT/IRIS \citep[][]{moore14,wright14}, JWST/NIRSPEC \citep[][]{closs08,birkmann14}, and JWST/MIRI-MRS \citep[][]{wells15} will trace the crucial rest-optical bandpass in galaxies out to redshift $z \sim 10$ and beyond. % Complications that the MaNGA pipeline must handle. Imaging spectroscopic surveys such as MaNGA face substantial calibration challenges in order to meet the science requirements of the survey \citep{yan16b}. In addition to requiring accurate absolute spectrophotometry from each fiber, MaNGA must correct for gravitationally-induced flexure variability in the Cassegrain-mounted BOSS spectrographs, determine accurate micron-precision astrometry for each IFU bundle, and combine spectra from the individual fibers with accurate astrometric information in order to construct 3-D data cubes that rectify the wavelength-dependent differential atmospheric refraction and (despite large interstitial gaps in the fiber bundles) consistently deliver high-quality {\it imaging} products. These combined requirements have driven a substantial software pipeline development effort throughout the early years of SDSS-IV. % \begin{deluxetable*}{lcccc} \tablecolumns{5} %\tablewidth{100pt} \tabletypesize{\scriptsize} \tablecaption{IFU Data Reduction Software} \tablehead{ \colhead{Telescope} & \colhead{Spectrograph} & \colhead{IFU} & \colhead{Pipeline} & \colhead{Reference}} \startdata \multicolumn{5}{c}{Fiber-Fed IFUs}\\ \hline AAT & AAOMEGA & SAMI & {\sc 2dfdr} & \citet{sharp15}\\ Calar Alto 3.5m & PMAS & PPAK & {\sc p3d} & \citet{sandin10} \\ & & & {\sc r3d} & \citet{sanchez06}\tablenotemark{a}\\ & & & IRAF & \citet{martin13}\tablenotemark{c}\\ HET & VIRUS & VIRUS & {\sc cure} & \citet{snigula14}\\ McDonald 2.7m & VIRUS-P & VIRUS-P &{\sc vaccine} & \citet{adams11}\\ & & & {\sc venga} & \citet{blanc13}\\ SDSS 2.5m & BOSS & MaNGA & {\sc mangadrp} & This paper\\ WHT & WYFFOS & INTEGRAL & {\sc iraf} & \\ WIYN & WIYN Bench Spec. & DensePak & {\sc iraf} & \citet{andersen06}\\ & & SparsePak & {\sc iraf} & \\ \hline \multicolumn{5}{c}{Fiber $+$ Lenslet-Based IFUs}\\ \hline AAT & AAOMEGA & SPIRAL & {\sc 2dfdr} & \citet{hopkins13}\\ Calar Alto 3.5m & PMAS & LARR & As PPAK above & \\ Gemini & GMOS & GMOS & {\sc iraf} & \\ Magellan & IMACS & IMACS & {\sc kungifu} & \citet{bolton07}\\ VLT & GIRAFFE & ARGUS & {\sc girbldrs} & \citet{blecha00} \\ & & & {\sc eso cpl}\tablenotemark{b} & \\ & VIMOS & VIMOS & {\sc vipgi} & \citet{zanichelli05}\\ & & & {\sc eso cpl}\tablenotemark{b} & \\ \hline \multicolumn{5}{c}{Lenslet-Based IFUs}\\ \hline Keck & OSIRIS & OSIRIS & {\sc osirisdrp} & \citet{krabbe04} \\ UH 2.2m & SNIFS & SNIFS & {\sc snurp} & \\ WHT & OASIS & OASIS & {\sc xoasis} & \\ & SAURON & SAURON & {\sc xsauron} & \citet{bacon01}\\ \hline \multicolumn{5}{c}{Slicer-Based IFUs}\\ \hline ANU & WiFeS & WiFeS & {\sc iraf} & \citet{dopita10}\\ Gemini & GNIRS & GNIRS & {\sc iraf} & \\ & NIFS & NIFS & {\sc iraf} & \\ VLT & KMOS & KMOS & {\sc eso cpl}\tablenotemark{b}, {\sc spark} & \citet{davies13}\\ & MUSE & MUSE & {\sc eso cpl}\tablenotemark{b} & \citet{weil12}\\ & SINFONI & SINFONI & {\sc eso cpl}\tablenotemark{b} & \citet{modig07}\\ \enddata \tablenotetext{a}{See \citet{sanchez12} for details of the implementation for the CALIFA survey.} \tablenotetext{b}{See http://www.eso.org/sci/software/cpl/} \tablenotetext{c}{Reference corresponds to the DiskMass survey.} \label{software.table} \end{deluxetable*} Historically, IFU data have been processed with a mixture of software tools ranging from custom built pipelines \citep[e.g.,][]{zanichelli05} to general purpose tools capable of performing all or part of the basic data reduction tasks for multiple IFUs. For fiber-fed IFUs (with or without coupled lenslet arrays) that deliver a pseudo-slit of discrete apertures the raw data is similar in format to traditional multiobject spectroscopy and has hence been able to build upon an existing code base. In contrast, slicer-based IFUs produce data in a format more akin to long-slit spectroscopy, while pure-lenslet IFUs are different altogether with individual spectra staggered across the detector. Following \citet[][]{sandin10}, we provide here a brief overview of some of the common tools for the reduction of data from optical and near-IR IFUs \citep[see also][]{bershady09}, including both fiber-fed IFUs with data formats similar to MaNGA and lenslet- and slicer-based IFUs by way of comparison. As shown in Table \ref{software.table}, the {\sc iraf} environment remains a common framework for the reduction of data from many facilities, especially Gemini, WIYN, and the WHT. Similarly, the various IFUs at the VLT can all be reduced with software from a common ISO C-based pipeline library, although some other packages \citep[e.g., GIRBLDRS,][]{blecha00} are also capable of reducing data from some VLT IFUs. Substantial effort has been invested in the {\sc p3d} \citep{sandin10} and {\sc r3d} \citep{sanchez06} packages as well, which together are capable of reducing data from a wide variety of fiber-fed instruments (including PPAK/LARR, VIRUS-P, SPIRAL, GMOS, VIMOS, INTEGRAL, and SparsePak) for which similar extraction and calibration algorithms are generally possible. For survey-style operations, the SAMI survey has adopted a two-stage approach, combining a general-purpose spectroscopic pipeline {\sc 2dfdr} \citep{hopkins13} with a custom three-dimensional stage to assemble IFU data cubes from individual fiber spectra \citep{sharp15}. Similarly, the MaNGA Data Reduction Pipeline ({\mangadrp}; hereafter the DRP) is also divided into two components. Like the {\sc kungifu} package \citep{bolton07}, the 2d stage of the DRP is based largely on the SDSS BOSS spectroscopic reduction pipeline {\sc idlspec2d} (Schlegel et al., in prep), and processes the raw CCD data to produce sky-subtracted, flux-calibrated spectra for each fiber. The 3d stage of the DRP is custom built for MaNGA, but adapts core algorithms from the CALIFA \citep{sanchez12} and VENGA \citep{blanc13} pipelines in order to produce astrometrically registered composite data cubes. In the present contribution, we describe version v1\_5\_4 of the MaNGA DRP corresponding to the first public release of science data products in SDSS Data-Release 13 (DR13)\footnote{DR13 is available at http://www.sdss.org/dr13/}. We start by providing a brief overview of the MaNGA hardware and operational strategy in \S \ref{hardwareops.sec}, and give an overview of the DRP and related systems in \S \ref{drpoverview.sec}. We then discuss the individual elements of the DRP in detail, starting with the basic spectral extraction technique (including detector preprocessing, fiber tracing, flatfield and wavelength calibration) in \S \ref{drp2d.1}. In \S \ref{skysub.sec} we discuss our method of subtracting the sky background (including the bright atmospheric OH features) from the science spectra, and demonstrate that we achieve nearly Poisson limited performance shortward of 8500 \AA. In \S \ref{fluxcal.sec} we discuss the method for spectrophotometric calibration of the MaNGA spectra, and in \S \ref{waverect.sec} our approach to resampling and combining all of the individual spectra onto a common wavelength solution. We describe the astrometric calibration in \S \ref{astrometry.sec}, combining a basic approach that takes into account fiber bundle metrology, differential atmospheric refraction, and other factors (\S \ref{basicast.sec}) and an `extended' astrometry module that registers the MaNGA spectra against SDSS-I broadband imaging (\S \ref{eam.sec}). Using this astrometric information we combine together individual fiber spectra into composite 3d data cubes in \S \ref{cubes.sec}. Finally, we assess the quality of the MaNGA DR13 data products in \S \ref{dq.sec}, focusing on the effective angular and spectral resolution, wavelength calibration accuracy, and typical depth of the MaNGA spectra compared to other extant surveys. We summarize our conclusions in \S \ref{summary.sec}. Additionally, we provide an Appendix \ref{datamodel.sec} in which we outline the structure of the MaNGA DR13 data products and quality-assessment bitmasks. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% % Hardware Design and Observing %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%

\label{summary.sec} The 13th data release of the Sloan Digital Sky Survey includes the raw MaNGA spectroscopic data, the fully reduced spectrophotometrically calibrated data, and the pipeline software and metadata required for individual users to re-reduce the data themselves. In this work, we have described the framework and algorithms of the MaNGA Data Reduction Pipeline software \mangadrp\ version v1\_5\_4 and the format and quality of the ensuing reduced data products. The DRP operates in two stages; the first stage performs optimal extraction, sky subtraction, and flux calibration of individual frames, while the second combines multiple frames together with astrometric information to create calibrated individual fiber spectra (in a row-stacked format) and rectified coadded data cubes for each target galaxy. The row-stacked spectra and coadded data cubes are provided for both a linear and a logarithmically sampled wavelength grid, both covering the wavelength range 3622 - 10,354 \AA. For the 1390 galaxy data cubes released in DR13 we demonstrate that the MaNGA data have nearly Poisson-limited sky subtraction shortward of $\sim$ 8500 \AA, with a residual pixel value distribution in all-sky test plates nearly consistent with a Gaussian distribution whose width is determined by the expected contributions from detector and Poisson noise. Each MaNGA exposure is flux calibrated independently of all other exposures using mini-bundles placed on spectrophotometric standard stars; based on comparison to broadband imaging the composite data cubes have a typical relative calibration of 1.7\% (between \Hb\ and \Ha) with an absolute calibration of better than 5\% for more than 89\% of the MaNGA wavelength range. These data cubes reach a typical $10\sigma$ limiting continuum surface brightness $\mu = 23.5$ AB arcsec$^{-2}$ in a five arcsec diameter aperture in the $g$ band. Additionally, we have demonstrated that: \begin{itemize} \item The wavelength calibration of the MaNGA data has an absolute accuracy of 5 \kms\ rms with a relative fiber-to-fiber accuracy of better than 1 \kms\ rms. \item The astrometric accuracy of the reconstructed MaNGA data cubes is typically 0.1 arcsec rms, based on comparison to previous SDSS broadband imaging. \item The spatial resolution of the MaNGA data is a function of the observational seeing, with a median of 2.54 arcsec FWHM. We have shown that the effective reconstructed point source profile is well described by a single gaussian whose parameters are given in the header of each data cube. \item The spectral resolution of the MaNGA data is a function of both both fiber number and wavelength, but has a median $\sigma = 72$ \kms. \end{itemize} Despite these overall successes of the MaNGA DRP, we conclude by noting that there is still ample room for future improvements to be made in some key areas. First, sky subtraction (while adequate for most purposes) shows some non-gaussianities in the residual distribution, a slight overstimate in the read noise of one camera, and a possible systematic oversubtraction at the $\sim 0.1\sigma$ level in the blue. Work is ongoing to test whether better treatment of amplifier crosstalk or the scattered light model can improve limiting performance in this area for the purposes of extremely deep spectral stacking. Secondly, the spectral line spread functions given in the DR13 data products (and in previous SDSS optical fiber spectra) are effectively under-reported by about 10\%. Work is currently underway to use high spectral resolution observations of MaNGA target galaxies to constrain this effect more precisely and fix it in future data releases. Third, spatial covariance in the reconstructed data cubes (treated here by a simple functional approximation) can also be treated more completely. Finally, with additional data it will be possible to fine tune the MaNGA quality control algorithms (which currently can be overly aggressive in flagging potentially problematic cases) and likely recover some of the objects whose reduced data have been identified as unreliable for use in DR13.

1607.08619

1607

1607.08157_arXiv.txt

{ Several directional techniques have been proposed for a directional detection of Dark matter, among others anisotropic crystal detectors, nuclear emulsion plates, and low-pressure gaseous TPCs. The key point is to get access to the initial direction of the nucleus recoiling due to the elastic scattering by a WIMP. In this article, we aim at estimating, for each method, how the information of the recoil track initial direction is preserved in different detector materials. We use the SRIM simulation code to emulate the motion of the first recoiling nucleus in each material. We propose the use of a new observable, \textit{D}, to quantify the preservation of the initial direction of the recoiling nucleus in the detector. We show that in an emulsion mix and an anisotropic crystal, the initial direction is lost very early, while in a typical TPC gas mix, the direction is well preserved. }

Weakly interacting massive particles are among the most studied candidates for Dark Matter (DM). The direct search for DM is performed by looking for elastic collisions of these WIMPs with target nuclei in a detector~\cite{goodman_1985}. For a 1\,GeV/c$^2$ to 1000\,GeV/c$^2$ DM particle, the kinetic energy of the recoiling nuclei would be in the range of few keV to few hundreds of keV depending on the mass of the target nuclei. The search for such rare events, at low energies, requires a highly efficient discrimination of the background with respect to signal composed of nuclear recoils. Non-directional detection is mainly limited -- after an ideal electron-recoil discrimination -- by two non-reducible background components: the neutrons and the neutrinos. Neutron background comes from natural radioactivity of the detector itself and the surrounding environment ($^{238}$U and $^{232}$Th) and from cosmic muons interacting with the rock and the detector. This background is present in all detectors, no matter how large the fiducial volume is. A veto can help to partly discriminate this background ; however the neutrons produced in the shielding or inside the detector are hard to discriminate. Neutrinos may come from the Sun, the atmosphere (interaction with cosmic rays) and from the DSNB (Diffuse Supernova Neutrino Background). The coherent scattering of these neutrinos on the target nuclei produces recoils in the energy region of interest (\mbox{E$_R$ $\lesssim$ 100\,keV}). When scaling the experiment, the neutrino background becomes non negligible. This ultimate neutrino background, called ``neutrino floor'', cannot be discriminated with standard non-directional methods, as no veto or shielding can be set for neutrinos and thus enforces a lower limit to the cross section attainable with non-directional detectors~\cite{billard_2014}. On the other hand, directional detection proposes to use the anisotropy in the recoil angular distribution (in the galactic coordinates) originated from the motion of the solar system around the galactic center within the DM halo~\cite{spergel_1988}. The mostly isotropic background can thus be distinguished from the (anisotropic) expected DM signal. In case of a claimed detection by a non-directional detector, a directional detection would be an unambiguous proof of the DM origin of the claimed signal~\cite{billard_2012}. The crucial point of the directional detection is to get access to the initial direction of the recoil nucleus. Several strategies have been proposed for a directional detection. We will focus on three strategies and present them: anisotropic crystals, nuclear emulsions and low pressure gaseous TPCs. The goal here is to provide a method to assess how the information of the initial direction of a nuclear recoil is retained in a given material. The direction actually measured by a detector is mainly defined by the WIMP-ion recoil angular distribution. Instrumental effects might dilute the initial direction information: readout \cite{battat_readout_2016}, diffusion effects. We will avoid the WIMP-ion kinematic effects, and will focus on the motion of the first recoiling nucleus in the detector material, and consider it has a fixed energy and a fixed direction. We use SRIM simulations~\cite{ziegler_2008,ziegler_2010} to emulate the motion of such nuclear recoil in each detector material; the SRIM outputs will help us define a figure of merit evaluating the preservation of the direction of a nuclear recoil in each material.

\label{sec:discussion} Using SRIM simulations, we compared how the recoil primary direction information is preserved in the detector materials for three different strategies for a Dark matter directional detection: anisotropic crystals, emulsion layers and low-pressure gaseous TPCs. Anisotropic crystals do not allow an event by event reconstruction of the tracks; yet showing how they preserve the direction information is useful to understand the prospective results of a directional measurement. On the other hand, emulsion and TPCs detectors measure each recoil track; the ranges of the typical recoil tracks expected from the elastic scattering by a WIMP is consistent with the reconstruction resolution of their respective readouts. We propose the use of a new observable, \textit{D}, to quantify the preservation of the initial direction information of the recoiling nucleus. This observable shows that among the three studied materials, a low pressure TPC gaseous mix is better suited for measuring the direction of WIMP-induced nuclear recoils. In fact, dedicated measurements with a calibration source producing ions with a known direction at a given kinetic energy such as the COMIMAC facility \cite{muraz_comimac_2016} would allow to confirm the simulations presented here. \vspace{5mm} We acknowledge F. Hosseini and Q. Riffard help during the phase of preliminary data analysis. CC acknowdleges financial support from the Labex Enigmass.

1607.08157

1607

1607.01847_arXiv.txt

We study the axionic inflation with a modulated potential and examine if the primordial tensor power spectrum exhibits oscillatory feature, which is testable with future space-based gravitational wave experiments such as DECIGO and BBO. In the case of the single-field axion monodromy inflation, it turns out that it is difficult to detect the oscillation in the spectrum due to suppression of the sub-Planckian decay constant of axion. On the other hand, in the case of aligned chromo-natural inflation where the axion is coupled to a SU(2) gauge field, it turns out that the sizable oscillation in the tensor spectrum can occur due to the enhancement of chiral gravitational waves sourced by the gauge field. We expect that this feature will be a new probe to axion phenomenologies in early universe through the chiral gravitational waves.

Inflation can be regarded as magnifier which allows us to probe microscopic world, namely, high energy physics through the cosmic microwave background (CMB) anisotropies or the large scale structures of the universe. Indeed, it is important to explore high energy physics such as string theory by means of inflation. It is known that there are a lot of scalar fields called axions and gauge fields in string theory. It should be noted that an axion is one of the best-motivated candidates of an inflaton since it naturally gives rise to a nearly flat potential protected by its shift symmetry \cite{Freese:1990rb}. One of the characteristic feature of axionic inflation is that it gets correction of the form of the periodic potential due to quantum non-perturbative effects such as instantons. Specific examples are the axion monodromy mechanism \cite{Silverstein:2008sg} or a kind of aligned natural inflation motivated by the weak gravity conjecture \cite{Kim:2004rp, delaFuente:2014aca}. Remarkably, this kind of potential gives the small modulation to the scalar power spectrum. Thus, the oscillatory feature in the power-spectrum is intimately related to fundamental physics. Intriguingly, although the oscillation in the scalar spectrum has been often discussed~\cite{Wang:2002hf, Easther:2013kla}, the oscillation in the tensor spectrum has been overlooked. This is because, on the CMB scales, the oscillation amplitude in the tensor spectrum is suppressed by several orders of magnitude compared to that in scalar spectrum. However, it is worth seeking the possibility of this oscillatory signature in the tensor spectrum on the scales probed by future space-based gravitational wave experiments such as DECIGO \cite{Kawamura:2011zz} and BBO \cite{Crowder:2005nr}. In this work, we explore the possibility of generating primordial gravitational waves (PGWs) with oscillatory features in axionic inflation. Specifically, we focus on two types of axionic inflations: one is single-field axion monodromy inflation and the other is a variant of inflation driven by the axion coupled to SU(2) gauge field, called aligned chromo-natural inflation \cite{Adshead:2012kp}. These two models are quite different from the point of the mechanism of producing tensor spectra. In the case of single-field monodromy inflation, the tensor spectrum comes from vacuum fluctuations. We see that it is difficult to detect the oscillatory feature from single-field monodromy inflation since the amplitude of oscillation is suppressed by the factor of slow-roll parameters and sub-Planckian decay constant. On the other hand, in the case of aligned chromo-natural inflation, the tachyonic growth of one helicity mode of the gauge field produces chiral PGWs during inflation \cite{Dimastrogiovanni:2012ew, Adshead:2013qp, Obata:2014loa, Obata:2016tmo}. We find that the tensor mode due to particle production of gauge field is sensitive to the modulation of inflaton potential and it produces detectable oscillatory feature in the tensor spectrum even for the tiny modulation. This feature will open up a new window to physics in early universe through the chiral gravitational waves.

We studied the oscillatory feature of tensor spectrum from axionic inflation. In the case of single-field monodromy inflation, the modulation in tensor spectrum is too small to be detected by DECIGO or BBO. On the other hand, in the case of aligned chromo-natural inflation, we can get the sizable modulation in one helicity mode of tensor perturbation sourced by the gauge field which experienced the tachyonic instability around horizon crossing during inflation. Thus, we found the possibility of producing sizable oscillatory feature in the tensor spectrum of chiral gravitational waves from axionic inflations when axion couples to the gauge field during inflation. In this work, we discussed the tensor spectrum with a sizable oscillation produced by non-Abelian gauge field in chromo-natural inflation. It is known that it is difficult to reconcile the original chromo-natural model with CMB data because it yields too large red scalar spectral index or too much chiral GWs. However, it is possible to improve the model so that chromo-natural inflation occurs in a frequency range higher than nHz and CMB constraints can be satisfied~\cite{Obata:2016tmo}. Moreover, we can expect the sizable modulation in the tensor spectrum in the case of Abelian gauge field because one helicity mode of gauge field produces tensor modes at the non-linear level \cite{Sorbo:2011rz}. We leave these issues for future work.

1607.01847

1607

1607.07699_arXiv.txt

Neutral atomic hydrogen (\schi) gas in interstellar space is largely organized into filaments, loops, and shells, the most prominent of which are ``supershells''. These gigantic structures requiring $\ga~3\times~10^{52}$~erg to form are generally thought to be produced by either the explosion of multiple supernovae (SNe) in OB associations or alternatively by the impact of high-velocity clouds (HVCs) falling to the Galactic disk. Here we report the detection of a kiloparsec (kpc)-size supershell in the outskirts of the Milky Way with the compact HVC 040+01$-$282 (hereafter CHVC040) at its geometrical center using the ``Inner-Galaxy Arecibo L-band Feed Array'' \schi\ 21-cm survey data. The morphological and physical properties of both objects suggest that CHVC040, which is either a fragment of a nearby disrupted galaxy or a cloud originated from an intergalactic accreting flow, collided with the disk $\sim 5$~Myrs ago to form the supershell. Our result shows that some compact HVCs can survive their trip through the Galactic halo and inject energy and momentum into the Milky Way disk.

\label{sec:intro} Supershells are large gaseous shells of radius greater than a few hundred parsecs. They are distinct from other shell-like structures in their extraordinarily large energy requirement, i.e., $\ga 3\times 10^{52}$~erg, which corresponds to $\ga 30$ supernova (SN) explosions \citep{heiles1979,heiles1984,mcclure2002}. About twenty supershells have been found in the Milky Way, and numerous neutral atomic hydrogen (\schi) holes corresponding to supershells have been discovered in nearby dwarfs and spiral galaxies \citep{kamphuis1991,bagetakos2011}. These gigantic structures are generally thought to arise from multiple SN explosions in stellar OB associations. But most supershells are missing a stellar association in their interior, and the number of supershells and their energies are usually incompatible with the level of star formation in those galaxies \citep{heiles1984,rhode1999}. Therefore, several alternative scenarios have been proposed, the most popular of which is the collision of high-velocity clouds (HVCs) with the disk \citep{tenorio1980,tenorio1987,mirabel1990}. HVCs are \schi\ clouds with radial velocities very different from the disk material in the Milky Way, e.g., with a deviation more than 50~\kms\ from the range of permitted velocities in a simple model of the distribution and rotation of the \schi\ gas in the Galaxy \citep{wakker1991a}. Some large HVC complexes are known to be gas streams tidally stripped from satellite galaxies of the Milky Way, but the origin of isolated compact HVCs (CHVCs) remain controversial: they could be clouds formed from galactic fountain or intergalactic accreting flows, part of the large HVC complexes, or condensations in the multi-phase circumgalactic medium \citep{putman2011, wakker2004, putman2012}. The HVC origin has been proposed for a few supershells \citep{heiles1984, mirabel1990,tamanaha1997}, but there has been no clear example showing a direct link between the two, particularly for CHVCs. Here we report the detection of a Galactic supershell with an associated HVC, GS040.2+00.6$-$70 (hereafter GS040). GS040 was first identified as a faint, forbidden-velocity wing feature (FVW~40.0+0.5) in the low-resolution, large-scale longitude-velocity study of \cite{kang2007}. We have found that GS040 appears to be a complete circular ring with complicated structures inside in our high-resolution I-GALFA (Inner-Galaxy ALFA) \schi\ 21-cm line survey data. The I-GALFA survey is a survey of the first Galactic quadrant visible to Arecibo ($\ell=32\arcdeg$ to 77\arcdeg\ and $|b| \la 15\arcdeg$) done by using the 7-beam Arecibo L-band Feed Array (ALFA) receiver on the Arecibo 305~m telescope, and it provides sensitive ($\Delta T_b=0.2$~K) and fully-sampled \schi\ maps at spatial and spectral resolutions of 4\arcmin\ and 0.184~\kms, respectively \citep{koo2010,gibson2012}. The I-GALFA survey data further reveal that there is a CHVC at the very center of GS040. This CHVC, named HVC 040+01$-$282 (hereafter CHVC040), was first identified in the Leiden/Dwingeloo survey \citep{wakker1991b} and was later classified as an isolated CHVC by \cite{braun1999} and \cite{deheij2002}. \citet{westmeier2005} presented a higher-resolution (9\arcmin) \schi\ image obtained from the Effelsberg telescope, which showed that CHVC040 has a pronounced head-tail structure. Our Arecibo \schi\ images reveal detailed spatial and velocity structure of CHVC040 strongly suggesting its association with the supershell GS040. We describe two structures in Section~\ref{sec:targets}, and discuss their physical characteristics and their association together with some implications on the disruption of HVCs in Section~\ref{sec:disc}.

\label{sec:disc} \subsection{Formation of GS040 by the Collision of CHVC040} \label{sec:disc_sub1} The location of CHVC040 at the geometrical center of GS040 suggests that their physical association is very likely. The centroids of the CHVC emission and that of the GS040's hub emission overlap within $\sim 0.2$~degrees. The probability of this being a random alignment is $\sim 3 \times 10^{-6}$, and multiplying this by $\sim 300$ CHVCs yields an overall probability of $9\times 10^{-4}$ for any CHVC aligning this well with the GS040's hub. No intermediate-velocity \schi\ connecting the two is apparent (Figure~\ref{fig:pvmap}), but this could be because the gas is ionized. We searched for a warm ionized gas associated with GS040 using the Wisconsin H$\alpha$ Mapper Northern Sky Survey (WHAM-NSS) data \citep{haffner2003}. The survey has an angular resolution of 1\arcdeg, and, around the GS040 area, it provides H$\alpha$ spectra covering $\vlsr$ from $-85$ to $+100$~\kms\ at spectral resolutions of 12~\kms. We have examined the H$\alpha$ intensity map integrated over the velocity range of GS040 ($-85$ to $-66$~\kms), but could not detect an associated emission ($\simlt$ 0.05~$R$ where $1 R = 10^6/4\pi\,{\rm photons\,cm}^{-2}\,{\rm sr}^{-1}\,{\rm s}^{-1}$). Note that this area is bright in the total H$\alpha$ intensity map with a mean intensity of about 3~$R$, so that we do not expect to see the faint emission associated with either GS040 or CHVC040 in the all-sky H$\alpha$ maps \citep{finkbeiner2003,dennison1998}. We also searched for an associated hot ionized gas using the 0.1-2.4 keV image of the ROSAT All-Sky X-ray Survey \citep[1 pixel scale = 44\arcsec;][]{voges1999}, but couldn't detect any emission. CHVC040 belongs to the ``Galactic Center Negative'' (GCN) HVC complex, which is a collection of small discrete HVCs sparsely distributed over a $70\arcdeg\times 70\arcdeg$ area within $\ell=0\arcdeg$ to $70\arcdeg$ and $b=-60\arcdeg$ to 10$\arcdeg$ \citep{wakker1991b,winkel2011}. A kinematic model has been proposed where GCN is a smooth gas flow starting at $b=-60\arcdeg$ at a heliocentric distance of 35~kpc and crossing the Galactic plane obliquely at 15~kpc \citep{jin2010}. For comparison, GS040 is probably located near the Scutum-Centaurus (Sct-Cen) arm at a distance of $\sim 20$~kpc \citep{dame2011} because the disrupted interstellar medium (ISM) is seen only at velocities below that of the Sct-Cen arm ($\sim -60$~\kms), not at higher velocities (see Figures~\ref{fig:chmap_gs040} and \ref{fig:pvmap}). We examined lists of known stellar objects, \schii\ regions, OB stars, and SN remnants (SNRs), but no known sources are likely to be associated with GS040. There is one \schii\ region (G$039.864+00.645$) in the {\it WISE} catalog of Galactic \schii\ regions \citep{anderson2014} that is located at $\sim 6\arcmin$ west of the hub. But this \schii\ region has a systematic velocity of $-40.9$~\kms\ \citep{anderson2011} and is enclosed by a small (68\arcsec) dust shell, so it cannot be responsible for the HI shell. Instead the distance of $\sim 20$~kpc is not unreasonable for CHVC040, because GCN does not have a smooth extended envelope like other HVC complexes, and it appears to be composed of several subpopulations that do not share a common origin \citep{winkel2011}. Note that the geometrical center of GS040 is well below the Galactic plane ($\sim 420\,\dtwenty$~pc where $\dtwenty\equiv d/20~{\rm kpc}$), i.e., at $b=+0\fdg6$ while the midplane there is at $b \sim +1\fdg8$ because the Galactic plane is warped in the outer Galaxy \citep{levine2006}. It is difficult to imagine the SN origin for a supershell at such height, and CHVC040 is most likely the energy/momentum source for GS040. The total energy deposited ($E_E$) in the Galactic disk by CHVC040 can be inferred from the parameters of the GS040 supershell. The radius of GS040 is $450\,\dtwenty$~pc, while its mass at $\vlsr \le -75$~\kms\ is $1.6\times 10^5\,\dtwenty^2$~\msol\ including the cosmic abundance of helium. If we account for the mass unobservable due to Galactic background \schi\ emission, the total mass of GS040 would be considerably greater. Adopting $v_s\sim 30$~\kms\ as the expansion speed of the shell, its kinetic energy is $E_K\ga 1.4\times 10^{51}\,\dtwenty^2$~erg. The collision should have occurred $\sim (1/3) R_s/v_s \sim 5\times 10^6\,\dtwenty$~yr ago, where the numerical factor $1/3$ accounts for the deceleration of the shell. Note that $E_K$ is a small fraction of the total energy deposited ($E_E$), most of which should have been radiated away. If GS040 was produced by multiple SNe, then, assuming instantaneous energy injection \citep{heiles1979}, $E_E\sim 5.3\times 10^{43}\, n_0^{1.12}\, R_s^{3.12}\, v_s^{1.4} \sim 1.2\times 10^{53}\, (n_0/0.1~{\rm cm}^{-3})$~erg, where $n_0$ is ambient hydrogen density. At the position of GS040, the mean \schi\ density in the midplane is $\sim 0.1\,{\rm cm}^{-3}$, and the \schi\ scale height is about 720~pc \citep{levine2006}. So $n_0\sim 0.06~{\rm cm}^{-3}$, and we have $E_E\sim 7\times 10^{52}$~erg. For comparison, the total extent of CHVC040 is $210\,\dtwenty\times 320\,\dtwenty\,{\rm pc}^2$ while its \schi\ mass is $\mchvc=5,\!800\,\dtwenty^2$~\msol. The area used to derive this mass has a geometrical mean radius of 150~pc, so that the mean hydrogen density of CHVC040 is 0.017~cm$^{-3}$. If CHVC040 collided with the rotating disk with the mean `deviation' speed (absolute difference in velocity from the disk gas there) of HVCs, i.e., $\sim 240$~\kms\ \citep{wakker2004}, the kinetic energy and momentum would be $4.7\times 10^{51}\,\dtwenty^2$~ergs and $2.0\times 10^{6}\,\dtwenty^2$~\msol\kms\ including the He abundance, respectively. This implies that the mass of CHVC040 when it collided with the disk should have been an order of magnitude greater and that the HVC that we see (CHVC040) might be the remains of the original HVC. The spatial morphology and velocity structure of CHVC040 suggest that it is moving southwest in the plane of the sky and is approaching us. The steep southwestern boundary of the cloud might represent the region compressed by the interaction with the ambient medium, whereas the diffuse envelope might be the material stripped off from the cloud due to the interaction or by the ram pressure of the surrounding medium \citep{santillan1999, kwak2009}. The small velocity width along the southwestern boundary and large velocity width beyond appear to be consistent with such speculation. Recently, \citet{heitsch2016} suggested that, for a CHVC with head-tail structure, the inclination angle, i.e., the angle between the CHVC's trajectory and the line-of-sight, can be derived from the asymmetry in position-velocity diagram along the head-tail line crossing the center of mass. The morphology of CHVC040 is different from typical head-tail CHVCs, i.e., CHVC040 has a wide flaring `tail' not a narrow elongated tail (see also \S~2), and its core is fragmented, so that it is not obvious if their model can be applied. Nevertheless if we assume that the center of mass is near the southwestern boundary and apply their model (their equations 1--4) to the position-velocity diagram along the red-dashed line in Figure 1, we obtain an inclination angle of $\sim 30\arcdeg$ which is consistent with our expectation. The shape of CHVC040 appears to be pointed away from the Galactic midplane as if it already went through the midplane. But the geometrical center of the supershell coincides with the current location of CHVC040, not being located in the midplane. Perhaps CHVC040 is approaching at an angle to the warped Galactic disk and has not yet fully penetrated the disk to the midplane. It is worth to note that, according to the \citet{jin2010}'s model, the GCN HVC stream is colliding to the Galactic plane almost perpendicularly {\em from below}, while the head-tail directions of individual HVCs seem to indicate that there is no preferential direction in their motions \citep{winkel2011}. Therefore, the orbit of CHVC040 is uncertain. The colliding geometry and the origin of the complex structures such as the hub and spokes in GS040 need to be explored. \subsection{CHVC040 and Disruption of HVCs} There are about three hundred known CHVCs in the Milky Way \citep{deheij2002, putman2002}. A considerable fraction of CHVCs has a head-tail structure indicating a ram pressure interaction with the diffuse galactic halo gas \citep{putman2011}. An important question is whether they are totally dissipated in the Galactic halo to feed the multi-phase circumgalactic medium or they can survive their trip through the halo \citep[e.g.,][]{putman2011}. Since CHVC040 is located in the far outer Galaxy, it may be of extragalactic origin rather than originating from a Galactic fountain, although it is not clear whether CHVC040 was originally a fragment of a nearby tidally disrupted galaxy or a cold cloud formed in a larger accreting flow of ionized, low-metallicity intergalactic gas. Our result then directly shows that at least some CHVCs of extragalactic origin do survive and collide with the Galactic disk. According to numerical studies, CHVCs with \schi\ masses $\la 3\times 10^4$~\msol\ would be totally disrupted in the Galactic halo unless they are embedded in dark matter \citep{heitsch2009, plockinger2012}. But dynamical shielding by an extended diffuse gaseous component can significantly extend their lifetime \citep{putman2012}. We have checked whether there are additional sources like the CHVC040-GS040 system in the I-GALFA \schi\ data using the HVC catalog of \cite{deheij2002}. There are twelve CHVCs in their Table~2 including CHVC040 in the I-GALFA survey area, most of which are at relatively high latitudes. They are isolated HVCs and have $\vlsr < -150$~\kms, so they belong to the GCN HVC complex. We see low-velocity \schi\ features around some CHVCs, but none of them appear associated. A systematic study against all CHVCs may reveal other CHVC-supershell systems.

1607.07699

1607

1607.02243_arXiv.txt

We present emission line templates for passively evolving (``retired") galaxies, useful for investigation of the evolution of the ISM in these galaxies, and characterization of their high-temperature source populations. The templates are based on high signal-to-noise ($>800$) co-added spectra ($3700-6800$\AA) of $\sim11500$ gas-rich Sloan Digital Sky Survey galaxies devoid of star-formation and active galactic nuclei. Stacked spectra are provided for the entire sample and sub-samples binned by mean stellar age. In Johansson~et al (2014), these spectra provided the first measurements of the He II 4686\AA\ line in passively-evolving galaxies, and the observed He II/H$\beta$ ratio constrained the contribution of accreting white dwarfs (the ``single-degenerate'' scenario) to the type Ia supernova rate. In this paper, the full range of unambiguously detected emission lines are presented. Comparison of the observed [O I] 6300\AA/H$\alpha$ ratio with photoionization models further constrains any high-temperature single-degenerate scenario for type Ia supernovae (with 1.5 $\lesssim$ T/$10^{5}K$ $\lesssim$ 10) to $\lesssim$3--6\% of the observed rate in the youngest age bin (i.e. highest SN Ia rate). Hence, for the same temperatures, in the presence of an ambient population of post-AGB stars, we exclude additional high-temperature sources with a combined ionizing luminosity of $\approx 1.35\times 10^{30} L_{\odot}/M_{\odot,*}$ for stellar populations with mean ages of 1 -- 4 Gyrs. Furthermore, we investigate the extinction affecting both the stellar and nebular continuum. The latter shows about five times higher values. This contradicts isotropically distributed dust and gas that renders similar extinction values for both cases.

1607.02243

1607

1607.02596_arXiv.txt

namefont}{\normalfont\bfseries} \usepackage{titlesec} % \titleformat{ \noindent New arguments supporting the reality of large-scale fluctuations in the density of the visible matter in deep galaxy surveys are presented. A statistical analysis of the radial distributions of galaxies in the COSMOS and HDF-N deep fields is presented. Independent spectral and photometric surveys exist for each field, carried out in different wavelength ranges and using different observing methods. Catalogs of photometric redshifts in the optical (COSMOS-Zphot) and infrared (UltraVISTA) were used for the COSMOS field in the redshift interval $0.1 < z < 3.5$, as well as the zCOSMOS (10kZ) spectroscopic survey and the XMM-COSMOS and ALHAMBRA-F4 photometric redshift surveys. The HDFN-Zphot and ALHAMBRA-F5 catalogs of photometric redshifts were used for the HDF-N field. The Pearson correlation coefficient for the fluctuations in the numbers of galaxies obtained for independent surveys of the same deep field reaches $R = 0.70 \pm 0.16$. The presence of this positive correlation supports the reality of fluctuations in the density of visible matter with sizes of up to 1\,000 Mpc and amplitudes of up to 20\% at redshifts $z \sim 2$. The absence of correlations between the fluctuations in different fields (the correlation coefficient between COSMOS and HDF-N is $R = -0.20 \pm 0.31$) testifies to the independence of structures visible in different directions on the celestial sphere. This also indicates an absence of any influence from universal systematic errors (such as "spectral voids"), which could imitate the detection of correlated structures. {\bf Key words:} Cosmology: observations; large-scale structure of Universe

Modern observational cosmology has led to the discovery of very large structures with scales of order 100 Mpc in the spatial distribution of galaxies in the local Universe, at redshifts $z \sim 0.1$, and also in the spatial distribution of quasars at redshifts $z \sim 2$ $[1-6]$. Over the last decade, observations of the largescale structure of the Universe [7] have moved from groups and clusters of galaxies with sizes of the order of 1 Mpc to structures with sizes of $\sim 100$ Mpc (SDSS superclusters, in particular, the Sloan Great Wall, with a size of 420 Mpc [1]). The mass distribution can be determined independently from analyses of the proper motions of galaxies. Tully et al. [2] recently discovered a coherent motion of galaxies forming the Laniakea (Local) Supercluster with a diameter of 160 Mpc. Groups of quasars with scales of $10-100$ Mpc have also been found, beginning with the study of Komberg et al. [3, $4-6$]. Modern multi-band photometric deep galaxy surveys can be used to study the spatial distribution of galaxies at redshifts $0.3-3$, which has led to the detection of inhomogeneities on scales up to 1\,000 Mpc $[8-10]$. A comparison of the wideangle Sloan Digital Sky Survey (SDSS) and the COSMOS pencil-beam survey is shown in Fig. 1, together with the radial distribution of the number of galaxies. Fluctuations in the numbers of galaxies in neighboring volume elements of a pencil-beam survey are due to the presence of Poisson noise (the discreteness of the sample), systematic observational errors (selection effects), and the presence of large-scale structure (the "cosmic variance"), which plays an important role in comparisons of models with the observations. The main difficulty in distinguishing real density fluctuations is the possibility of hidden selection effects that are present in each galaxy survey, which can imitate large-scale inhomogeneity of the galaxy distribution. \begin{figure*} \centering \includegraphics[scale=0.35,clip]{Shirokov_fig1} \hfill \includegraphics[scale=0.35,clip]{Shirokov_fig2} \label{SDSS} \caption{ Upper: distribution of galaxies in the SDSS. The one-square-degree COSMOS deep field is marked by the dark strip. Lower: fit of the radial distribution of the number of galaxies using a uniform distribution. The gray area highlights regions where there are deficits or excesses of galaxies.} \end{figure*} In the current study, we present new arguments supporting the reality of large-scale fluctuations in the matter density in deep galaxy surveys.

Our analysis of the radial distribution of galaxies in the COSMOS/UVISTA field shows that the real observed fluctuations in the spatial distribution of the galaxies appreciably exceed the predictions of the $\Lambda$CDM model for the evolution of non-baryonic dark matter, in both their amplitude and linear size. This means that the $\Lambda$CDM model requires the introduction of a large bias factor ($b\sim10$ relative to the non-baryonic dark matter) at redshifts $z\sim1$. It is also necessary to explain the large scale of the positive correlation corresponding to the linear size of the detected structures. This follows from Table 3, which shows that the fluctuations in the number of galaxies preserve their sign over several adjacent bins, while neighboring bins should have opposite signs in the $\Lambda$CDM model. Thus, in addition to the difficulties of the $\Lambda$CDM model on small scales (galaxies and halos with sizes of $10-100$ kpc [28, 29]), there also exist problems on very large scales, associated with the presence of large-scale inhomogeneities in the spatial distribution of galaxies with sizes of the order of 1\,500 Mpc and amplitudes exceeding 20\%. \begin{figure} \centering \includegraphics[scale=0.66]{Shirokov_fig13} \caption{Comparison of $\delta_{obs}$ for the COSMOS (solid curve) and HDF-N (dotted curve) catalogs. The Poisson noise shown (dot-dashed curve) corresponds to $\sigma_{P}$ for the HDF-N catalog.} \label{HDF-N_COSMOS} \end{figure} Our analysis of the COSMOS and UVISTA survey data and comparisons of these data with the data from the ALHAMBRA, zCOSMOS, and XMMCOSMOS surveys leads to the following conclusions. \begin{itemize} \item The detected inhomogeneities in the radial distribution of galaxies in the COSMOS, UVISTA, ALHAMBRA, and XMMphot-COSMOS photometric catalogs is confirmed by data from the zCOSMOS and XMMspec-COSMOS spectroscopic catalogs, and these data are mutually consistent. \item The amplitudes and linear sizes of the fluctuations in the independent COSMOS (optical) and UltraVISTA (near infrared) catalogs, and also in the ALHAMBRA/Field 4, XMMNewton and zCOSMOS catalogs, are mutually consistent. The corresponding correlation coefficient is positive and equal to $\rho > 0.5$. \item The amplitudes and sizes of the fluctuations are stable for different fits and various limiting redshifts $z_{max}$. When the bin size is decreased, the fluctuation amplitude grows, and the individual density peaks coincide with galaxy clusters detected earlier. \end{itemize} The amplitudes and sizes we have found agree with the amplitudes and sizes of inhomogeneities found for the COSMOS field in other studies using the 10k-zCOSMOS spectral survey [22, 30], ALHAMBRA photometric survey [21], and X-ray observations [31]. Appreciable fluctuations in the number density of galaxies in slices of the COSMOS survey at various redshifts were found in [32], where is it emphasized that these structures really exist. Thirty-six candidate structures were found at redshifts $1.5 < z < 3.1$, having masses of $10^{15}M_{\odot}$. The sizes of the observed radial structures appreciably exceed the transverse cross sections of the pencil-beam surveys. The detected individual galaxy clusters fall near peaks of the fluctuations found using small redshift-bin sizes [25]. In particular, the peak at $z = 0.73$ corresponds to a galaxy cluster that was detected in [26] using spectroscopic observations, and the three peaks at $z\sim0.35$, $z\sim0.7$, and $z\sim0.85$ coincide with clusters detected in [22], as well as in our own study. The paper [21] describes the ALHAMBRA catalog, which includes the region of the COSMOS survey. This catalog also displays non-uniformity in the radial distribution of galaxies, which is correlated with the non-uniformity observed for the COSMOS survey. X-ray sources from the COSMOS catalog are considered in [31]. Peaks in the radial distribution of these X-ray sources agree with regions where there are excess galaxies in the optical and IR, indirectly supporting the presence of large-scale structures. The detection of fluctuations in the number of galaxies, manifest in the same way in independent observations and obtained using independent datareduction methods, substantially reduces the possibility that these are associated with unknown systematic errors. This suggests with a high degree of certainty that the fluctuations observed in the COSMOS/UVISTA field are related to the cosmic variance, and thus imply positive correlations in the spatial distribution of galaxies in deep surveys. \subsection*

1607.02596

1607

1607.00841_arXiv.txt

H{\sc i} absorption studies of active galaxies enable us to probe their circumnuclear regions and the general interstellar medium, and study the supply of gas which may trigger the nuclear activity. In this paper, we investigate the detection rate of H{\sc i} absorption on the nature of radio galaxies based on their emission-line spectra, nature of the host galaxies based on the \textit{WISE} colours and their radio structure, which may help understand the different accretion modes. We find significant difference in distributions of W2$-$W3 colour for sources with H{\sc i} absorption detections and non-detections. We report a high detection rate of H{\sc i} absorption in the galaxies with \textit{WISE} infrared colours W2$-$W3 $>$ 2, which is typical of gas-rich systems, along with a compact radio structure. The H{\sc i} detection rate for low-excitation radio galaxies (LERGs) with W2$-$W3 $>$ 2 and compact radio structure is high (70.6$\pm$20.4 \%). In HERGs, compact radio structure in the nuclear or circumnuclear region could give rise to absorption by gas in the dusty torus in addition to gas in the interstellar medium. However, higher specific star formation rate (sSFR) for the LERGs with W2$-$W3 $>$ 2 suggests that H{\sc i} absorption may be largely due to star-forming gas in their hosts. LERGs with extended radio structure tend to have significantly lower values of W2$-$W3 compared to those with compact structure. Extended radio sources and those with W2$-$W3 $<$ 2 have low H{\sc i} detection rates.

Optical spectroscopic studies of radio galaxies \citep{1979MNRAS.188..111H,1994ASPC...54..201L, 2010A&A...509A...6B,2012MNRAS.421.1569B} led to the division of radio-loud AGNs (Active Galactic Nuclei) into two categories, HERGs (High Excitation Radio Galaxies) and LERGs (Low Excitation Radio Galaxies). This dichotomy is believed to be due to differences in the accretion modes of HERGs and LERGs \citep{2010A&A...509A...6B, 2012MNRAS.421.1569B, 2012ApJ...757..140S}. In the high-excitation or quasar mode or radiatively efficient mode of accretion, Eddington ratio is greater than 1\%, while it is less than 1\% in the low-excitation mode or radio mode or radiatively inefficient mode of accretion \citep{2014ARA&A..52..589H}. In HERGs, accretion in the nuclear regions takes place through geometrically thin, optically thick accretion disks \citep{1973A&A....24..337S,1973blho.conf..343N} while for LERGs, this thin disk is absent and hence the radiatively inefficient accretion mode is believed to operate \citep{1994ApJ...428L..13N, 1995ApJ...452..710N, 2014ARA&A..52..529Y}. HERGs or sources with higher accretion rate are galaxies with younger stellar populations, growing central black hole masses and higher rate of central star formation \citep{2008MNRAS.384..953K,2012MNRAS.421.1569B}. Contrary to this, LERGs or sources with lower excitation and lower accretion rate are galaxies of higher stellar masses, higher black hole masses and with redder optical colours, consisting older stellar population than HERGs \citep{2008MNRAS.384..953K,2012MNRAS.421.1569B}. While LERGs contain both FR\,I and FR\,II type sources \citep{1974MNRAS.167P..31F}, HERGs are predominantly FR\,II \citep{2012MNRAS.421.1569B, 2014ARA&A..52..589H}. Differences are also seen in their luminosities. A majority of LERGs have a radio luminosity less than $\sim$ 10$^{26}$ W Hz$^{-1}$ at 1.4 GHz, while HERGs dominate above this luminosity \citep{2012MNRAS.421.1569B}. It has been suggested that different kinds of fuelling processes may be dominant in different kinds of AGN \citep{2004IAUS..222..235M}. While interactions and major mergers are thought to be behind the fuelling of HERGs \citep{2014MNRAS.445L..51T, 2015ApJ...806..147C}, for low luminosity radio-loud AGNs, which would include most of the LERGs, accretion of hot halo ISM or IGM gas or minor mergers have been suggested as possible fuelling mechanisms by a number of authors \citep{2006MNRAS.372...21A, 2007MNRAS.376.1849H, 2007NewAR..51..168B, 2008A&A...486..119B, 2015MNRAS.451L..35E}. However, although many theoretical models are based on the framework where HERGs are fuelled by major mergers, the observational evidence is mixed, and different processes may be operating \citep{2014ARA&A..52..589H}. In the case of LERGs, major mergers are unlikely to be the triggering and fuelling mechanism \citep{2015MNRAS.451L..35E}. Recent studies \citep{2014MNRAS.444.3408Y} of early-type galaxies find presence of cold atomic/molecular gas in $\sim$40 \% of galaxies in the nearby Universe, and there is also evidence that early-type galaxies with dust lanes have higher chances to host emission-line AGNs \citep{2012MNRAS.423...59S}. \cite{2015MNRAS.449.3503D} argue that early-type galaxies with dust lanes have acquired their dust and gas from recent external minor mergers. For sources in clusters of galaxies, accretion of hot halo gas which has cooled also provides a viable mechanism \citep{2014ARA&A..52..589H}. In their infra-red study of a sample of southern 2 Jy radio galaxies, \cite{2014MNRAS.445L..51T} found host galaxies to have enough cold ISM fuel needed for the central AGN, but they argue that only having enough fuel is not sufficient for triggering the AGN activity, and it also depends on the kinematics and distribution of cold gas. H{\sc i} absorption studies towards radio-loud AGNs have been a prominent tool to study cold neutral hydrogen gas in the central regions of the host galaxies of these sources \citep{1989AJ.....97..708V, 2001MNRAS.323..331M, 2003A&A...404..871P, 2003A&A...404..861V, 2006MNRAS.373..972G,2010MNRAS.406..987E, 2011MNRAS.418.1787C, 2012MNRAS.423.2601A, 2013MNRAS.429.2380C, 2015A&A...575A..44G}. Studying the H{\sc i} properties of these sources with different radio luminosities, optical and infrared properties may provide clues towards further understanding and distinguishing between the accretion modes. However, most of the H{\sc i} absorption studies have been predominantly towards higher radio luminosity sources (L$_\mathrm{1.4 {GHz}}$ $>$ 10$^{26}$ WHz$^{-1}$), due to senstivity limitations of the instruments while observing flux density limited samples, except in a few cases e.g. \citep{2010MNRAS.406..987E,2011MNRAS.418.1787C, 2014A&A...569A..35G, 2015A&A...575A..44G}, where luminosities have been lower than 10$^{26}$ WHz$^{-1}$. In our studies \citep{2011MNRAS.418.1787C}, we observed the Compact Radio sources at Low Redshift (CORALZ) core sample of 18 sources, which was compiled by \cite{2004MNRAS.348..227S}, with the sources having flux densities larger than 100 mJy at 1400 MHz and angular sizes less than 2 arcsec. With a small sample and large statistical uncertainties, H{\sc i} properties were found to be similar to those in higher luminosity sources. A recent study by \cite{2014A&A...569A..35G, 2015A&A...575A..44G} of a much larger sample of 101 sources has H{\sc i} absorption detected towards 32 sources. In this paper, we explore possible dependences of H{\sc i} absorption properties of the sources on HERGs/LERGs and hence on the accretion mode, source size, and nature of the host galaxy as reflected by the infrared colours. We consider the sample observed uniformly by \cite{2015A&A...575A..44G} and the classification into high-excitation and low-excitation radio galaxies done by \cite{2012MNRAS.421.1569B} for a large sample of sources using the SDSS optical spectroscopic data (DR7; \citealt{2009ApJS..182..543A}) and the \emph{Faint Images of the Radio Sky at Twenty Centimeters survey} (FIRST; \citealt{1995ApJ...450..559B}) and \emph{NRAO VLA Sky Survey} (NVSS; \citealt{1998AJ....115.1693C}) radio data. We also used the\emph{ Wide-field Infrared Survey Explorer} (\emph{\textit{WISE}}; \citealt{2010AJ....140.1868W}) infra-red (IR) archival data to look for relation between IR and H{\sc i} gas properties of LERGs and HERGs. \textit{WISE} data provide mid-infrared photometry at different wavelengths (W1: 3.4 $\mu$m, W2: 4.6 $\mu$m, W3: 12 $\mu$m, W4: 22 $\mu$m). \textit{WISE} colour plots are useful to distinguish between optical AGN activity (W1$-$W2; \citealt{2013MNRAS.436.3451S}) and emission from warm dust heated due to AGN or star formation (W2$-$W3). W2$-$W3 colour can be also used as an indicator of star formation history in a galaxy \citep{2012ApJ...748...80D}, and hence helpful in distinguishing IR \lq early' type and IR \lq late' type galaxies \citep{2010AJ....140.1868W}. We explore how the H{\sc i} absorption gas properties change with W2$-$W3 colour for LERGs and HERGs. \begin{figure*} \centering \hbox{ \includegraphics[scale=0.44]{Fig1left_Sep1.eps} \includegraphics[scale=0.44]{Fig1right_Sep1.eps} } \caption{ Left: O[{\sc iii}] equivalent width vs. excitation index (for 91 sources with all 6 emission lines) with filled symbols (detections) and empty symbols(non-detections), vertical line represents EI=0.95 while horizontal line is for O[{\sc iii}] equivalent width = 5 \AA; Right: Log O[{\sc iii}]/H$\beta$ vs. Log N[{\sc ii}]/H$\alpha$ (for 99 sources with atleast 4 emission lines), solid green curve is the \protect\cite{2001ApJ...556..121K,2006MNRAS.372..961K} dividing line between AGNs and composite (SF+AGN) galaxies while dashed blue curve is the \protect\cite{2003MNRAS.341...33K} dividing line between star-forming and composite galaxies. } \label{fig1} \end{figure*} \begin{figure*} \centering \hbox{ \includegraphics[scale=0.44]{Fig2topleft_Sep2.eps} \includegraphics[scale=0.44]{Fig2topright_Sep1.eps} } \hbox{ \includegraphics[scale=0.44]{Fig2bottomleft_Sep2.eps} \includegraphics[scale=0.44]{Fig2bottomright_Sep2.eps} } \caption{ Top Left: Distributions of W2$-$W3 colours for H{\sc i} absorption detections (blue) and non-detections (red); Top Right: W1$-$W2 vs. W2$-$W3 plot for detections (filled symbols) and non-detections (empty symbols) for the different categories of objects. Horizontal red dashed line shows W1$-$W2 $=$ 0.8 and vertical red dashed line shows W2$-$W3 $=$ 2 Bottom: Distributions of W2$-$W3 colours for H{\sc i} absorption detections and non-detections shown separately for LERGs (left) and HERGs (right). } \label{fig2} \end{figure*}

We summarise our conclusions in this section. \begin{enumerate} \item Earlier studies of H{\sc i} absorption had reported that compact radio sources, namely the compact steep-spectrum and giga-Hertz peaked spectrum sources, exhibited the highest detection rates of up to $\sim$45 per cent. We have shown that there is significant difference in distributions of W2$-$W3 colour for H{\sc i} absorption detections and non-detections. Considering the galaxies with \emph{WISE} infrared colour W2$-$W3 $>$ 2, which is typical of gas-rich systems, along with a compact radio structure leads to high detection rates of over 70 per cent. \item Although majority of LERGs have low H{\sc i} detection rate, the compact LERGs with bright \emph{WISE} colours W2$-$W3 $>$ 2 also have high H{\sc i} detection rates of over 70 per cent, which indicates a gas and dust rich ISM. \item The distributions of W2$-$W3 for compact and extended LERGs are significantly different, with the extended LERGs having lower values, the maximum difference in their cumulative distributions occurring at W2$-$W3 $=$ 1.8. The W2$-$W3 colours appear to be playing an important role in determining detection of H{\sc i} in absorption rather than the specific star-formation rate. \item Overall a lower rate of H{\sc i} detection in LERGs is consistent with a scenario of suppressed star formation rate suggested in the literature, perhaps due to feedback from radio source. It is likely that the LERGs are following a secular evolutionary process and some of these started their evolution with higher gas/dust content and may have also undergone a minor merger at some stages of their evolution, but feedback from the radio source appears to be playing a role in affecting and suppressing star formation at later stages. \end{enumerate}

1607.00841

1607

1607.05713_arXiv.txt

The Disk Detective citizen science project aims to find new stars with 22 $\mu$m excess emission from circumstellar dust using data from NASA's WISE mission. Initial cuts on the AllWISE catalog provide an input catalog of 277,686 sources. Volunteers then view images of each source online in 10 different bands to identify false-positives (galaxies, background stars, interstellar matter, image artifacts, etc.). Sources that survive this online vetting are followed up with spectroscopy on the FLWO Tillinghast telescope. This approach should allow us to unleash the full potential of WISE for finding new debris disks and protoplanetary disks. We announce a first list of 37 new disk candidates discovered by the project, and we describe our vetting and follow-up process. One of these systems appears to contain the first debris disk discovered around a star with a white dwarf companion: HD 74389. We also report four newly discovered classical Be stars (HD 6612, HD 7406, HD 164137, and HD 218546) and a new detection of 22 $\mu$m excess around a previously known debris disk host star HD 22128.

\label{sec:introduction} All-sky mid-infrared surveys have revolutionized the science of planet formation by discovering populations of young stars and main sequence stars with excess infrared radiation indicating the presence of dusty circumstellar disks. These disks, which include gas-rich protoplanetary disks around Young Stellar Objects (YSOs) and dusty debris disks around main sequence stars, serve as the signposts of planet formation \citep[e.g.,][]{2002ApJ...577L..35K}. They inform us about the timescales and the environment of planet formation \citep[e.g.,][]{2005ApJ...620.1010R, 2007ApJ...671.1784H, 2015ApJ...808..167J}, and the present day locations and dynamics of planets \citep[e.g.,][]{2010ApJ...718L..87T, 2012ApJ...748L..22M, 2013ApJ...766L...1Q, 2015ApJ...798...83N, 2015ApJ...807L...7C}. The IRAS all sky survey discovered the first extrasolar debris disks \citep{1984ApJ...278L..23A} and provided a large sample of debris disks \citep[e.g.,][]{2007ApJ...660.1556R}. After IRAS, AKARI surveyed the whole sky at 9 and 18 $\mu$m with $\sim 7$ times better sensitivity than IRAS, finding many more new disks \citep{2010A&A...514A...1I}. Some disk discoveries have come from pointed studies, like the Spitzer Formation and Evolution of Planetary Systems (FEPS) survey \citep{2009ApJS..181..197C}. But many of the best-studied, most informative disks (like TW Hydra, Fomalhaut, etc.) are relatively isolated on the sky, requiring an all-sky survey to find them. NASA's Wide-field Infrared Survey Explorer (WISE) is the most recent and sensitive all-sky mid-infrared survey \citep{2010AJ....140.1868W}, with a further factor of $\sim 80$ gain in sensitivity over AKARI in the mid-IR. Using a 16-inch mirror in a Sun-synchronous polar orbit, WISE scanned the sky at 3.4 $\mu$m, 4.6 $\mu$m, 12 $\mu$m, and 22 $\mu$m (bands W1, W2, W3, and W4 respectively). The WISE cryogenic mission, launched in 2009, lasted a little over 10 months and was followed by the first post-cyrogenic mission, NEOWISE. The AllWISE catalog\footnote{Available at http://irsa.ipac.caltech.edu/Missions/wise.html} combines data from both phases, making it the most comprehensive mid-infrared multi-epoch view of the sky available today. Previous infrared surveys for debris disks have provided target lists for exoplanet searches via direct imaging \citep{2008ApJ...672.1196A, 2013ApJ...773...73J, 2013ApJ...773..179W, 2015ApJ...800....5M}. Debris disks found with WISE should provide crucial targets for upcoming generations of exoplanet searches. WISE could detect debris disks around main sequence A stars to a distance of 300 pc and protoplanetary disks around T Tauri stars to 1 kpc. Indeed, many teams have used the WISE data to find new debris disks, searching a vast catalog of $>747$ million WISE sources. \citet{2012MNRAS.427..343M} cross-correlated the WISE source list with the Hipparcos catalog, finding over 86,000 stars with suspected infrared excesses. \citet{2013MNRAS.433.2334K}, \citet{2013ApJS..208...29W} and \citet{2014ApJS..212...10P} performed more careful searches for debris disks in the WISE source list using the Hipparcos catalog and found 6, 70 and 108 new debris disk candidates, respectively. Other specific surveys for debris disks have focused on stars with ages determined from chromospheric activity \citet{2014ApJ...780..154V}, white dwarfs \citep{2011ApJ...729....4D, 2012ApJ...759...37D}, M dwarfs \citep{2012A&A...548A.105A, 2013AAS...22115813O}, G-K dwarfs \citep{2014MNRAS.437..391C}, Kepler candidate exoplanet systems \citep{2012ApJ...752...53L, 2012A&A...541A..38R, 2012MNRAS.426...91K} and other exoplanet catalogs \citep{2012ApJ...757....7M}. Likewise, the WISE data on young clusters and star-forming regions have attracted much attention. \citet{2011ApJS..196....4R, 2014ApJ...784..126E, 2014AJ....147..133L} scoured the Taurus-Auriga Region. \citet{2012ApJ...744..130K} searched 11 outer Galaxy massive star-forming regions and three open clusters. Other studies have examined smaller regions, like the Western Circinus molecular cloud \citep{2011ApJ...733L...2L}, the young open cluster IC 1805 \citep{2013A&A...554A...3S} the \ion{H}{2} region S155 \citep{2014RAA....14.1269H}, the Sco-Cen and $\eta$ Cha associations \citep{2012MNRAS.421L..97R, 2012ApJ...758...31L}, nearby moving groups of young stars \citep{2012ApJ...751..114S} and $\lambda$ and $\sigma$ Orionis \citep{2015AJ....150..100K}. Still others have attempted to take in the whole sky, using color cuts \citep{2013Ap&SS.344..175M} or cross correlating with IRAS \citep{2014ApJ...784..111L}. Many of these searches were based on preliminary data releases with less sensitivity than the AllWISE release, but they have already uncovered thousands of candidate Class I, II and III YSOs and transitional disks, helping fill in our picture of the timing and progression of star formation. Unfortunately, because of its limited spatial resolution (12 arcsec at 22 $\mu$m) contamination and confusion limit every search for disks with WISE \citep[e.g.][]{2012MNRAS.426...91K}. Contamination sources include unresolved companion stars and other stars nearby on the sky, background galaxies, Galactic cirrus, and even asteroids and airplanes. For this reason, most recent searches include visual inspection of the WISE images \citep[see e.g.,][]{2011ApJ...729....4D, 2013ApJS..208...29W, 2014MNRAS.437..391C, 2014ApJS..212...10P}. Computer cuts alone can provide a first stage of vetting, but they generate catalogs riddled with false positives \citep{2012MNRAS.426...91K}; . Color cuts and source quality flags can help \citep[e.g.,][]{2012ApJ...744..130K, 2013Ap&SS.344..175M, 2014MNRAS.440.3430D}, but the color loci of disk candidates overlaps with the color loci of blended background galaxies \citep{2012ApJ...744..130K} and peaks in the Galactic dust emission. \citet{2012MNRAS.426...91K} used the IRAS 100 $\mu$m level to discard many false positive disk candidates contaminated with Galactic dust emission, but using this method prohibits searching many interesting star-forming regions. Because of these challenges, many disks remain to be found with WISE data, even after all the efforts described above. The largest published study of debris disks \citep{2013ApJS..208...29W} and the still larger WISE science team disk study (Padgett et al. in prep.) are based on the Hipparcos and Tycho catalogs. These catalogs are magnitude limited in V band, so they omit a vast population of redder, late type stars\footnote{The initial candidates listed in this paper are also all in Hipparcos, but our full dataset does not cross-correlate with any stellar catalogues.}. Moreover, a vast solid angle in young clusters and star-forming regions remains to be properly searched with WISE---each candidate examined by eye and followed up with spectroscopy and higher resolution imaging. When \citet{2013Ap&SS.344..175M} ran the all-sky data through a novel color filter to search for YSO candidates (without vetting the candidates by eye), he found a total of $\sim 10,000$ objects of interest; the WISE studies of young clusters and star-forming regions described above (which mostly included visual vetting) yielded a total of $\sim 4000$ disk candidates. The difference between these numbers provides a minimal measure of what remains for us to study with WISE: $\gtrsim 6000$ objects with colors consistent with YSOs that have not yet been visually inspected. Here we describe a new project to scour the WISE data for new debris disks and YSOs. The Disk Detective citizen science/crowdsourcing project classifies WISE sources via a website, diskdetective.org, where volunteers examine images from WISE, the Two Micron All Sky Survey (2MASS), the Digitized Sky Survey (DSS) and when available, the Sloan Digital Sky Survey (SDSS), to check them for false positives. This approach should allow us to unleash the full potential of WISE for finding new disks, probing the cooler stars and isolated objects missed by previous debris disk searches, a catalog 8 times the size of the large \citet{2013ApJS..208...29W} survey. We describe the online vetting process in Section \ref{sec:approach}, our small-telescope follow-up program in Section \ref{sec:followup} and we present our first list of 37 disk candidates in Section \ref{sec:candidates}.

\label{sec:conclusions} We have outlined and demonstrated a novel process for identifying new candidate circumstellar disks in the WISE survey data. This paper reports only results from the first 10\% of the search, so it might be premature to try to derive any statistically meaningful inferences about the population of debris disk from this limited sample. But our list of 37 new, well-vetted disk candidates demonstrates the utility of crowdsourcing analysis of WISE images. One of our disk candidate systems appears to contain the first debris disk discovered around a star with a white dwarf companion: HD 74389. We also report four newly discovered classical Be stars (HD 6612, HD 7406, HD 164137, and HD 218546) and a new detection of 22 $\mu$m excess around previously known debris disk host star HD 22128. We decided to only publish in this paper candidates that are in the Hipparcos catalog. Since the \citet{2013ApJS..208...29W} cross-correlated the WISE archive with the Hipparcos catalog, they could conceivably have identified all of the candidates that we are announcing. However, \citet{2013ApJS..208...29W} used a [2MASS] - [W4] color criterion, as opposed to our [W1] - [W4] color criterion, and used the WISE AllSky data rather than the WISE ALLWISE data. Yet all of the candidates presented in this paper would have also been selected by the \citet{2013ApJS..208...29W} color criterion. More importantly, while we examined candidates by eye to discard objects potentially contaminated by nearby stars and galaxies, \citet{2013ApJS..208...29W} used a statistical likelihood-ratio (LR) technique to accomplish this goal. Perhaps their statistical technique was more conservative than our more labor intensive approach, leaving these candidates unidentified. We have several further improvements to Disk Detective project underway, which we will describe in upcoming papers, including: \begin{itemize} \item New ways to retire the sources after fewer classifications. \item Spectroscopy of Southern hemisphere DDOIs via the CASLEO telescope in Argentina \item Imaging follow-up of DDOIs with the Robo-AO \citep{2014ApJ...790L...8B} instrument at the Palomar Observatory 60-inch telescope to check for background contaminants located closer to the star than DSS can probe. \end{itemize} With its high sensitivity and angular resolution in the mid-infrared, we expect that the James Webb Space Telescope (JWST) will be an important tool for following up disks discovered via Disk Detective. So we aim to have the project mostly completed by the time JWST launches in the fall of 2018.

1607.05713

1607

1607.07534_arXiv.txt

{We present deep SINFONI {\it K} band integral field spectra of two submillimeter (SMG) galaxy systems: BR 1202-0725 and J1000+0234, at $z=4.69$ and $4.55$ respectively. Spectra extracted for each object in the two systems do not show any signature of the [O\,{\sc ii}]$\lambda\lambda$3726,29\AA\,emission-lines, placing upper flux limits of $3.9$ and $2.5 \times 10^{-18}\,$\ergscm for BR 1202-0725 and J1000+0234, respectively. Using the relation between the star formation rate (SFR) and the luminosity of the [O\,{\sc ii}] doublet from \citet{kennicutt98}, we estimate unobscured SFR upper limits of $\sim$ $10-15\,$\myr and $\sim$ $30-40\,$\myr for the objects of the two systems, respectively. For the SMGs, these values are at least two orders of magnitude lower than those derived from SED and IR luminosities. The differences on the SFR values would correspond to internal extinction of, at least, $3.4 - 4.9$ and $2.1 - 3.6$ mag in the visual for BR 1202-0725 and J1000+0234 SMGs, respectively. The upper limit for the [O\,{\sc ii}]-derived SFR in one of the LAEs (Ly$\alpha2$) in the BR1202-0725 system is at least one order of magnitude lower than the previous SFR derived from infrared tracers, while both estimates are in good agreement for Ly$\alpha$1. The lower limits to the internal extinction in these two Lyman-alpha emitters (LAEs) are $0.6$ mag and $1.3$ mag, respectively. No evidence for the previously claimed \citep{ohta00} [O\,{\sc ii}] emission associated with Ly$\alpha$1 is identified in our data, implying that residuals of the K-band sky emission lines after subtraction in medium-band imaging data could provide the adequate flux.}

\begin{figure*} \centering \includegraphics[width=\textwidth]{figs/HST_image.eps} \caption{Left panel: {\it HST-NICMOS 2} image of the BR1202-0725 system, using the F160W filter, covering the SINFONI FoV of $8\arcsec \times 8\arcsec$. White crosses mark the positions of the SMG, quasar, Ly$\alpha1$ and Ly$\alpha2$ as labeled, according to \citet{carilli13}. White squares show the regions where spectra were extracted from the SINFONI datacube. Right panel: {\it HST-WFC3} image of J1000+0234, using the IR channel and F160W filter. White squares show the regions where SMG spectra were extracted from the SINFONI datacube, as identified. The white arrow points to a foreground object.} \label{hst} \end{figure*} Recent studies of galaxies detected at millimeter and submillimeter wavelengths have largely increased our understanding of the formation and evolution of massive galaxies when the Universe was 1-3 Gyrs old. Cosmological simulations indicate that massive galaxies can form at high-$z$ via gas-rich mergers, triggering extreme events such as intense star formation and simultaneously growing of supermassive black holes (SMBHs) close to their maximum (Eddington) accretion rates \citep{li07,netzer14}. Submillimeter galaxies (SMGs) play a key role in this scenario, since they represent examples of starburst galaxies (above main sequence) in the distant universe. At such high redshifts ($z\sim5$), a large fraction of the star formation activity is enshrouded in dust, most of the bolometric luminosity is radiated into the far-infrared (FIR, $40-500\,\mu$m) and submillimeter wavelengths and therefore the luminosity of galaxies is proportional to the star formation rate (SFR). SMGs observed with {\it Herschel} and ALMA \citep{mor12,vieira13,netzer14} have displayed extremely high infrared luminosities ($L_{FIR} \geq 10^{13} L_\odot$), which implies SFRs of $> 10^3 M_\odot\,$year$^{-1}$. These rates indicate that the bulk of the star formation in these galaxies have timescales of only $\leqslant 100\,$Myr. The structure and physical mechanisms at work in these objects are largely unknown. Different scenarios have been proposed that imply the removal of the available gas due to an active galactic nucleus (AGN) driven wind in luminous quasars or to supernovae and stellar winds in extreme starbursts, leaving behind a compact remnant \citep{sanders88,hopkins08,wuyts10}. However, evidence for significant outflow rates at high-$z$ are very limited and have only recently been achieved. One example is the highly magnified galaxy at $z = 4.92$ behind the lensing cluster MS 1358+62 \citep{swinbank09}, in which deep [O\,{\sc ii}]$\lambda\lambda$3726,29\AA\, spectroscopy indicates the presence of a young outflow ($< 15\,$Myr) and a SFR of $42$ \myr. In this work, we present deep SINFONI--IFS data of two $z \sim 5$ SMGs, BR 1202-0725 and COSMOS J100054+023436. BR 1202-0725 ($z = 4.69$) is a system composed by a dusty, luminous starburst (the SMG itself), an optically luminous QSO and two Ly$\alpha$ emitting extented regions, hereafter Ly$\alpha1$ and Ly$\alpha2$ \citep{omont96,carilli02}. The QSO and the SMG are separated by $3\farcs8$, with Ly$\alpha1$ between them ($2\farcs3\,$ north-west of the QSO), a clear signature of interaction, while Ly$\alpha2$ lies $2\farcs7\,$ south-west of the QSO (see left panel of Figure \ref{hst}). Recent observations of the system using ALMA detects narrow [C\,{\sc ii}]$\lambda158\,\mu$m emission in all four sources \citep{wagg12,carilli13}. Both SMG and QSO have high FIR luminosity ($L_{FIR} > 10^{13} L_\odot$) indicating SFRs of the order of $10^3\,$\myr or above \citep{iono06,carniani13}, while both Ly$\alpha$ emitting regions, Ly$\alpha1$ and Ly$\alpha2$, appear to be forming stars, added together, at a rate of $19$ \myr \citep{williams14}, derived from the [C\,{\sc ii}] $158 \mu$m emission-line luminosities. \citet{carilli13} concluded that the proximity of a luminous quasar is unlikely to be the source of the ionized nebula in Ly$\alpha1$, and speculate that a optically thick torus is shielding the radiation towards the Lyman-alpha emitters (LAEs, i.e. Ly$\alpha$1 and Ly$\alpha$2). Gas outflows related to the QSO are under debate, with some evidence based on the detection of a broad [C\,{\sc ii}] secondary component \citep{carilli13}, not confirmed by similar subsequent studies \citep{carniani13}. \citet{ohta00} report detection of [O\,{\sc ii}] doublet emission of the Ly$\alpha1\,$ galaxy based on narrow-band imaging. COSMOS J100054+023436 (hereafter J1000+0234, $z = 4.55$) is a SMG dominated by a starburst forming stars at a rate of $> 1000\,$\myr with a young age of $2-8\,$ Myr, estimated using infrared and radio measurements \citep{capak08}. This SMG presents multiple components observed in the Ly$\alpha$ emission-line, distributed along a region of $\approx 3''$ in size, and residing at the same redshift. This galaxy presents different morphology at different wavelengths, from the UV (Ly$\alpha$) to the near-infrared (rest-frame optical with [O\,{\sc ii}] doublet), likely indicating the presence of spatially resolved stellar populations on arcsec scales \citep{capak08}. Detection of a broad CO($4-3$) emission line indicates the presence of a large amount of molecular gas with an estimated total mass of $2.6 \times 10^{10}\,\textrm{M}_\odot$. The width of the CO($4-3$) line and the highly disturbed morphology suggest this system is involved in an ongoing merger \citep{schinnerer08}. The objective of this work is to search the [O\,{\sc ii}]$\lambda\lambda$3726,29\AA\, doublet in these two SMGs, to study the structure of the un-obscured star formation, including the presence of flows of gas other than rotation. This paper is organized as follows: in Sec. \ref{obs} we describe the observations and data reduction, in Sec. \ref{res} we present the results from the SINFONI datacubes for both SMG systems. In Sec. \ref{analys} we discuss the implications for the derived star formation rate, internal extinction in the visual and the claim of previous [O\,{\sc ii}] detection in the Ly$\alpha1$ of the BR1202-0725 system. Finally, conclusions are summarized in Sec. \ref{conc}. Thoughout this paper we assume a standard concordance cosmology ($H_0 = 70$, $\Omega_M = 0.3$, $\Omega_\Lambda = 0.7$).

\label{conc} In this work we presented deep SINFONI {\it K}-band integral field spectroscopy of two SMG systems at $z \sim 5$, BR1202-0725 and J1000+0234. BR1202-0725 consists of the SMG, a QSO and two LAEs while J1000+0234 appears to have a complex extended structure as traced by multiband optical and near-infrared imaging. Our main conclusions are: \begin{enumerate} \item The spectra extracted for all the objects, including the two LAEs, do not show any signature of the [O\,{\sc ii}] doublet, within the 5$\sigma$ sky uncertainties. The corresponding upper limits for the SFR in the two SMGs are two orders of magnitude below those derived from far-infrared measurements. The differences are explained as due to internal obscuration equivalent to an average visual internal extinction of, at least, $4.1$ mag and $2.8$ mag for BR1202-0725 and J1000+0234, respectively. These average high internal extinctions are similar to those measured in low-z U/LIRGs; \item The SFR upper limit for Ly$\alpha2$ derived from the non-detection of the [O\,{\sc ii}] is at least one order of magnitude lower than that measured from the infrared, while the corresponding values for Ly$\alpha1$ are in fairly good agreement, suggesting a very low internal extinction in this LAE. We find an internal extinction of at least $0.6$ and $1.3$ magnitudes in the visual for Ly$\alpha1$ and Ly$\alpha2$, respectively; \item Previous claims of [O\,{\sc ii}] emission associated to BR1202-0725 Ly$\alpha1$ based on narrow-band imaging are not confirmed by our SINFONI 2D spectroscopy. Residuals due to some of the sky emission lines within the filter bandpass could produce fluxes compatible with the [O\,{\sc ii}] flux reported previously by \citet{ohta00}. \end{enumerate}

1607.07534

1607

1607.07969_arXiv.txt

text{ The calibration of Very Long Baseline Interferometry (VLBI) data has long been a time consuming process. The Korean VLBI Network (KVN) is a simple array consisting of three identical antennas. Because four frequencies are observed simultaneously, phase solutions can be transferred from lower frequencies to higher frequencies in order to improve phase coherence and hence sensitivity at higher frequencies. Due to the homogeneous nature of the array, the KVN is also well suited for automatic calibration. In this paper we describe the automatic calibration of single-polarisation KVN data using the KVN Pipeline and comparing the results against VLBI data that has been manually reduced. We find that the pipelined data using phase transfer produces better results than a manually reduced dataset not using the phase transfer. Additionally we compared the pipeline results with a manually reduced phase-transferred dataset and found the results to be identical. } \begin{document} \jkashead %

} The calibration and analysis of Very Long Baseline Interferometry (VLBI) data has long been a time consuming process. However, with increases in computing power and hardware, it is becoming increasingly feasible to largely automate the process. The Korean VLBI Network (KVN) is an homogeneous radio array comprising of three antennas capable of simultaneously observing at 14\,mm, 7\,mm, 3.5\,mm and 2.3\,mm. The homogeneity of the array and the simultaneous multi-frequency capabilities allow for the transfer of phase solutions from lower frequencies to higher frequencies, make the KVN a suitable network to implement the automatic calibration of VLBI data. Arrays such as the KVN can produce as much as four datasets simultaneously, creating a large workload for those reducing the data. Due to this and other factors, there is a large backlog of non-reduced data, not just on the KVN but on most VLBI networks. For this reason, there is a large need for either fully automating or at least partially automatising the reduction process. The technical means for producing a fully automatic pipeline for single-polarisation continuum data are already developed with the largest constraints being typically logistical. A station should ideally provide calibration information (e.g. system temperature, antenna temperature and gain information) and also a monitoring log of the telescope, VLBI recorder, VLBI back-end and receiver. This information should allow for data to be automatically flagged. In many current VLBI experiments this information is unavailable or difficult to acquire. The KVN however largely provides this information. In this paper we describe a fully automatic pipeline for the KVN. The European VLBI Network (EVN) is a large consortium of observatories that collaborate to perform regular VLBI experiments \citep{evn}. Two pipelines have been developed to aid the calibration of EVN data. Before 2006, a pipeline was written as a special procedure within the Astronomical Image Processing System (AIPS) \citep{aips}. The pipeline broadly followed the calibration procedure described in Section \ref{sec:calibration}, but was not completely automated - although it greatly speed up the calibration process \citep{evnpipe1}. A more recent pipeline also implemented in AIPS is the task \emph{VLBARUN}, which applied phase and amplitude calibration to Very Long Baseline Array (VLBA) data before preceding to produce a simple image. A newer pipeline was developed for the automated reduction of EVN data, written in the Python AIPS wrapper, ParselTongue and provides largely similar functionality to the earlier pipeline except while being more robust and easier to use \citep{parsel}. This paper describes the automatic reduction software (pipeline) of an example KVN dataset and compares the results of the KVN Pipeline with that of the same dataset reduced in the manual way. In Section \ref{sec:observations}, we describe the observations that were used to test the KVN Pipeline, in Section \ref{sec:calibration}, we describe the traditional approach to VLBI calibration, in Section \ref{sec:KVNpipe} we describe the operation of the KVN Pipeline, in Section \ref{sec:discussion} we discuss the results and compare them against the traditionally reduced VLBI data and in Section \ref{sec:con} we present our conclusions.

} The KVN pipeline has demonstrated its ability to automatically reduce a single polarisation KVN datset without direct human intervention. The performance of the KVN Pipeline is consistently superiour to manually reduced data without FPT being applied. While the KVN is a simple network consisting of 3 identical telescopes, the methods described here (although not including the FPT capability) should be applicable to less homogoneous arrays such as the GMVA, KaVA and the EAVN. Currently the pipeline is only capable of reducing single polarisation continuum data. In future versions of the pipeline, we wish to add a polarisation reduction capability, spectral line reduction and possibly automated imaging and phase referencing capabilities.

1607.07969

1607

1607.06368_arXiv.txt

Outbursts on young stars are usually interpreted as accretion bursts caused by instabilities in the disk or the star-disk connection. However, some protostellar outbursts may not fit into this framework. In this paper, we analyze optical and near-infrared spectra and photometry to characterize the 2015 outburst of the probable young star ASASSN-15qi. The $\sim 3.5$ mag brightening in the $V$ band was sudden, with an unresolved rise time of less than one day. The outburst decayed exponentially by 1 mag for 6 days and then gradually back to the pre-outburst level after 200 days. The outburst is dominated by emission from $\sim10,000$ K gas. An explosive release of energy accelerated matter from the star in all directions, seen in a spectacular cool, spherical wind with a maximum velocity of 1000 \kms. The wind and hot gas both disappeared as the outburst faded and the source returned to its quiescent F-star spectrum. Nebulosity near the star brightened with a delay of 10--20 days. Fluorescent excitation of H$_2$ is detected in emission from vibrational levels as high as $v=11$, also with a possible time delay in flux increase. The mid-infrared spectral energy distribution does not indicate the presence of warm dust emission, although the optical photospheric absorption and CO overtone emission could be related to a gaseous disk. Archival photometry reveals a prior outburst in 1976. Although we speculate about possible causes for this outburst, none of the explanations are compelling.

Large brightness changes of young stars were seen long before the class of objects was understood to be related to star formation, or even that stars formed \citep{hind,ceraski06}. The largest and most prominent physical changes are decades-long bursts of accretion at rates of $10^{-4}-10^{-5}$ M$_\odot$ yr$^{-1}$, called FUor objects \citep[after FU Ori; see reviews by][]{herbig77,hartmann96,reipurth10}, and months-long bursts of accretion at $\sim10^{-7}$ M$_\odot$ yr$^{-1}$ \citep{aspin10,lorenzetti12}, called EXor outbursts \citep[after EX Lup;][]{herbig89}. The FUor and EXor classes of outbursts form our framework for interpreting large luminosity increases of young stars \citep[see review by][]{hartmann16}. The different timescales for the longer (FUor) and shorter (EXor) outbursts suggest that they are different phenomena. FUor outbursts are thought to be triggered by instabilities in the disk \citep[e.g.,][]{armitage01,vorobyov05,zhu09,bae14} at radii with rotational timescales of decades. On the other hand, the shorter EXor outbursts are thought to be triggered by instabilities in the magnetic connection between the star and disk \citep{dangelo10,dangelo12}. These two phenomena require different disk structures, which is supported with spectroscopic evidence. Viscously heated, optically thick FUor disks produce low-gravity spectra similar to those of supergiant stars, including deep absorption in CO overtone bands at 2.3 $\mu$m \citep[e.g.,][]{greene08}. EXors show a forest of optical emission lines, evidence of magnetospheric accretion, and a warm disk surface layer that produces strong CO emission \citep[e.g.,][]{herbig89,aspin10,caratti13,holoien14,sicilia15,banzatti15}. Since measuring the duration of an outburst usually requires impatient people to wait, these spectroscopic proxies are used to immediately discriminate between FUor and EXor outbursts. In practice, outbursts of protostars are often forced into the EXor/FUor classification scheme, even when the outburst does not fit well into either category. Many outbursts appear to be intermediate between the FUor and EXor classes \citep[e.g.,][]{contreras16}, despite the very different mechanisms and masses thought to be involved. In other cases, some characteristics of the outburst are inconsistent with this classification scheme \citep{ninan15}. Either the classification system groups together diverse physics, or the same accretion burst physics of EXors produces a wider range of observed phenomena than expected. Powerful new transient surveys are now leading to the discovery of 1--2 large outbursts of young stellar objects (YSOs) each year \citep[e.g.,][]{miller11,holoien14}. A recent outburst of a candidate young star, ASASSN-15qi, was identified by the All-Sky Automated Survey for Supernovae (ASAS-SN) variability survey \citep{shappee14}. Follow-up observations were obtained because the object was suspected to be young \citep[see discussion in][]{hillenbrand15} based on its projected spatial location near the H II regions Sh 2-148 and Sh 2-149 and the small, low extinction molecular cloud TGU 676/Dobashi 3359 \citep{dobashi05,dobashi11}. This region is a few degrees southwest of the main Cepheus star forming complexes, including LDN 1218 \citep[e.g.,][]{kun08,allen12}. A distance of $3.24$ kpc is adopted from the parallax measured for G108.47$-$2.81, which is visually located 1.5 deg from ASASSN-15qi and shares the same radial velocity \citep{choi14}. Initial optical spectroscopy of ASASSN-15qi showed P Cygni profiles characteristic of strong winds, an indicator of accretion-ejection processes that occur on young stars \citep{maehara15,hillenbrand15,connelley15}. However, subsequent low resolution near-infrared (IR) and high resolution optical spectra revealed a confusing picture of an outburst that does not neatly fit into either the FUor or EXor category \citep{hillenbrand15,connelley15}. Spatially extended emission in optical imaging \citep{hillenbrand15,stecklum15b} confirmed that the object is likely young. The source quickly faded from its peak \citep{stecklum15a,stecklum15b}. In this paper, we analyze optical and near-IR photometry and spectroscopy of the 2015 outburst of ASASSN-15qi. The outburst is especially remarkable because of (1) the dramatic 3.5 mag brightening in less than one day, and (2) the presence of a fast wind that either caused or was produced by the outburst. The wind faded as the outburst decayed over 3 months. In \S 2, we describe the wide array of observations used in this paper. In \S 3, we analyze these data to arrive at some empirical conclusions. In \S 4, we summarize the source properties, compare these properties to those of known outbursts where the physics is better understood, and describe some possible alternatives. These descriptions and speculations are summarized in \S 5.

1607.06368

1607

1607.01062_arXiv.txt

The Seyfert 1 galaxy, Ark 120, is a prototype example of the so-called class of bare nucleus AGN, whereby there is no known evidence for the presence of ionized gas along the direct line of sight. Here deep ($>400$\,ks exposure), high resolution X-ray spectroscopy of Ark\,120 is presented, from {\it XMM-Newton} observations which were carried out in March 2014, together with simultaneous {\it Chandra}/HETG exposures. The high resolution spectra confirmed the lack of intrinsic absorbing gas associated with Ark\,120, with the only X-ray absorption present originating from the ISM of our own Galaxy, with a possible slight enhancement of the Oxygen abundance required with respect to the expected ISM values in the Solar neighbourhood. However, the presence of several soft X-ray emission lines are revealed for the first time in the {\it XMM-Newton} RGS spectrum, associated to the AGN and arising from the He and H-like ions of N, O, Ne and Mg. The He-like line profiles of N, O and Ne appear velocity broadened, with typical FWHM widths of $\sim5000$\,km\,s$^{-1}$, whereas the H-like profiles are unresolved. From the clean measurement of the He-like triplets, we deduce that the broad lines arise from gas of density $n_{\rm e}\sim10^{11}$\,cm$^{-3}$, while the photoionization calculations infer that the emitting gas covers at least 10\% of $4\pi$ steradian. Thus the broad soft X-ray profiles appear coincident with an X-ray component of the optical--UV Broad Line Region on sub-pc scales, whereas the narrow profiles originate on larger pc scales, perhaps coincident with the AGN Narrow Line Region. The observations show that Ark\,120 is not intrinsically bare and substantial X-ray emitting gas exists out of our direct line of sight towards this AGN.

Photo-ionised or ``warm'' absorbers are commonly observed in at least 50\% of the UV/X-ray spectra of Seyfert 1s and type-1 QSO and are an important constituent of AGN (Reynolds 1997, Crenshaw, Kraemer \& George 2003, Porquet et al.\ 2004, Blustin et al. 2005, McKernan et al. 2007, Turner \& Miller 2009). Indeed the Seyfert warm absorbers that are frequently observed at high spectral resolution with {\sl XMM-Newton} and {\sl Chandra} are now known to give rise to numerous narrow absorption lines, usually blue-shifted -- implying outflowing winds -- of a few hundred km\,s$^{-1}$ up to a few thousand km\,s$^{-1}$. These arise from various elements over a wide range of ionisation levels and column densities, especially from iron, oxygen, carbon, nitrogen, neon, silicon and sulfur (Kaastra et al. 2000, Kaspi et al., 2002, McKernan et al. 2003, Blustin et al. 2003). Signatures range from the lowly ionised Unresolved Transition Array (UTA) of M-shell iron ($<$ Fe\,\textsc{xvii}) at $\sim16-17$\AA~ (Sako et al. 2001, Behar et al. 2001), to absorption from highly ionised (H-like and He-like) iron which may originate from an accretion disk wind (Reeves et al. 2004, Risaliti et al. 2005a, Braito et al. 2007, Turner et al., 2008, Tombesi et al. 2010, Gofford et al. 2013). These spectroscopic measurements can reveal crucial information on the outflow kinematics, physical conditions and locations relative to the central continuum source -- ranging from the inner nucleus (0.01 pc) to the galactic disk or halo (10 kpc). However a small class of nearby Seyfert galaxies exist which show no (or very little) X-ray or UV absorption. These AGN are the so-called {\it ``bare nucleus''} Seyferts or bare AGN. In principle the lack of intrinsic absorbing gas in these bare AGN allows a clean measurement of the innermost regions of the AGN and of the central engine closest to the black hole, removing any uncertainties as to how the absorbing gas is modeled. Ark\,120 (or Arakelian\,120) is a nearby ($z=0.032713$, Osterbrock \& Phillips 1977, Theureau et al. 2005) and X-ray bright ($F_{0.5-10\,{\rm keV}}=5.3\times10^{-11}$\,erg\,cm$^{-2}$\,s$^{-1}$, $F_{14-195\,{\rm keV}}=7.0\times10^{-11}$\,erg\,cm$^{-2}$\,s$^{-1}$; Patrick et al. 2011, Baumgartner et al. 2013) bare nucleus Seyfert 1. Along with its sister AGN, Fairall 9 (Emmanoulopoulos et al. 2011, Lohfink et al. 2012), it is the prototype example of a bare AGN. Indeed it is one of the brightest and cleanest bare AGN known, displaying neither intrinsic reddening in its IR/optical continuum nor evidence for absorption in UV and X-rays (Crenshaw et al.\ 1999, Reynolds 1997), allowing a clear view of the innermost regions of the AGN. A further key advantage for studying Ark\,120 is that it has a well determined reverberation based black hole mass, of $M_{\rm BH}$=1.5$\times$10$^{8}$\,M$_{\odot}$ (Peterson et al. 2004). An open question is whether Ark\,120 is intrinsically bare and devoid of circumnuclear X-ray emitting and/or absorbing gas, which may pose a challenge for unified schemes of AGN that imply the existence of wide scale obscuring and emitting gas (Antonucci 1993). Indeed, Vaughan et al. (2004) presented an initial 100\,ks {\it XMM-Newton} observation of Ark\,120 in 2003, which from the spectra obtained with the RGS (Reflection Grating Spectrometer, den Herder et al. 2001) showed no significant soft X-ray emission or absorption features associated to the AGN. Furthermore the X-ray continuum was found to be smooth from the soft X-ray band up to 10 keV, with a large but featureless soft X-ray excess present at energies below 2\,keV. This was also confirmed in a {\it Suzaku} study by Nardini et al. (2011), who favored a relativistic accretion disk reflection origin for the soft X-ray excess and broad iron K$\alpha$ line. In an alternative explanation for the broad band spectrum, Tatum et al. (2012) accounted for the iron K$\alpha$ emission through Compton scattering off an accretion disk wind, which had to be viewed out of the direct line of sight in this AGN. Most recently, Matt et al. (2014) presented a simultaneous $\sim 100$\,ks {\it XMM-Newton} and {\it NuSTAR} observation of Ark\,120 obtained in 2013 and showed that the bare broad-band X-ray spectrum could be explained by Comptonization of UV photons through a warm scattering medium associated to an accretion disk corona. This paper is the first of a series of papers to report upon the analysis of an unprecedented deep observational campaign on Ark\,120, which was subsequently obtained with {\it XMM-Newton} in 2014, with a total exposure exceeding 400\,ks (PI, D. Porquet). Part of the long {\it XMM-Newton} observations were performed simultaneously with {\it NuSTAR} to provide broad band hard X-ray coverage and with the High Energy Transmission Grating (HETG, Canizares et al. 2005) on-board {\it Chandra} to provide a high resolution view of the iron K band region. Here we concentrate on the high signal to noise and high resolution soft X-ray spectrum obtained with the RGS spectrometer on-board {\it XMM-Newton}. The primary goal is to determine whether the soft X-ray spectrum of Ark\,120 is intrinsically bare and devoid of circumnuclear X-ray gas, or indeed whether there are any signatures of ionized emission or absorption, which could arise from the accretion disk, the AGN broad and narrow line regions or from a nuclear outflow. We also present a search for any soft X-ray emission lines at high resolution from the Chandra/HETG above 1\,keV from Mg, Si, S. Subsequent papers will report in detail on the modeling of the iron K$\alpha$ line profile obtained as part of these observations (paper II, Nardini et al. 2016), as well as the nature of the broad band UV to hard X-ray continuum of Ark\,120 (paper III, Porquet et al. 2016, in preparation). The paper is organized as follows. In Section\,2, we describe the analysis of the RGS observations, while in Section\,3 the overall properties of the soft X-ray RGS spectrum are presented. Section\,4 is devoted to the analysis of the Galactic ISM absorption towards Ark\,120 and the subsequent modeling of the soft X-ray continuum. Section\,5 then describes the first detection of soft X-ray emission lines from Ark\,120 that have now been made possible through the deep RGS exposure and Section\,6 discusses their potential origin in the broad and narrow lined regions from the AGN. Values of H$_{\rm 0}=70$\,km\,s$^{-1}$\,Mpc$^{-1}$, and $\Omega_{\Lambda_{\rm 0}}=0.73$ are assumed throughout and errors are quoted at 90\% confidence ($\Delta\chi^{2}=2.7$), for 1 parameter of interest. All spectral parameters are quoted in the rest-frame of the AGN, at $z=0.032713$, unless otherwise stated. A conversion between energy and wavelength of $E = (12.3984$ \AA/$\lambda)$\,keV is adopted throughout.

\label{sec:discussion} \subsection{The Origins of the Ionized Soft X-ray Emission in Ark\,120} While there is no intrinsic X-ray absorption towards Ark 120 (aside from the neutral ISM absorption due to the Milkyway), the high signal to noise RGS spectrum has revealed several broad and narrow emission line profiles associated to the AGN. As shown in Section\,5, these broad profiles are associated with the He-like emission from N\,\textsc{vi}, O\,\textsc{vii} and Ne\,\textsc{ix} with velocity widths in the range from 4000--8000\,km\,s$^{-1}$ (FWHM). In addition, the line profile modeling in Section\,5.3 appears to exclude the possibility of the He-like triplets being composed of purely a blend of narrow unresolved emission lines from the forbidden, intercombination and resonance components, with the intrinsic velocity widths being resolved by the RGS (see Figure 8). In contrast the H-like profiles from N\,\textsc{vii}, O\,\textsc{viii} and Mg\,\textsc{xii} are narrow, for instance for the O\,\textsc{viii} Ly$\alpha$ emission the upper limit on the Gaussian velocity width is $\sigma<470$\,km\,s$^{-1}$. Thus while there is no direct absorption along the line of sight associated to the AGN (see also Vaughan et al. 2004, Matt et al. 2014), the fact that there is significant soft X-ray line emission associated with Ark\,120 suggests the AGN is not intrinsically bare. Thus the lack of absorption may just indicate that we are viewing the AGN along a preferential line of sight, with relatively little ionized gas along our direct view. The detection of narrow soft X-ray emission from many Seyfert galaxies, associated with photoionized or photoexcited gas, has proven common from grating observations of obscured AGN (Kinkhabwala et al. 2002), with the origin likely to be on scales consistent with the AGN Narrow Line Region (NLR) and perhaps arising from a large scale outflow. However recent observations are now also revealing the detections of broad soft X-ray line profiles from several Seyfert 1 galaxies, with velocity widths of several thousand km/s, suggesting an origin commensurate with the AGN Broad Line Region (BLR). Examples include NGC\,4051 (Ogle et al. 2004, Pounds \& Vaughan 2011), Mrk\,335 (Longinotti et al. 2008), Mrk\,841 (Longinotti et al. 2010), 3C\,445 (Reeves et al. 2010), Mrk\,509 (Detmers et al. 2011), MR\,2251-178 (Reeves et al. 2013) and NGC\,5548 (Kaastra et al. 2014). Indeed another AGN that bares some similarity to Ark\,120 is the bare Seyfert 1 galaxy, Mrk\,590. Here \xmm\ and \chandra\ showed no intrinsic absorption, but the presence of narrow ionized emission associated with highly ionized iron (Fe\,\textsc{xxv} and Fe\,\textsc{xxvi}) as well as from ionized Oxygen (O\,\textsc{viii}); see Longinotti et al. (2007). In addition, the extended soft X-ray emission detected in the {\it Chandra} image of Mrk\,590 implies the presence of ionized gas on larger (kpc) scales. The overall picture may be similar to what is observed in Ark\,120, namely that little ionized gas is detected along the line of sight, but evidence for photoionized gas is still seen from the circumnuclear gas out of the direct view. In Ark\,120, the detection of both narrow and broad lines implies the existence of soft X-ray emitting gas over a wide range of spatial scales, from the sub-pc BLR gas out to the more distant NLR, with the narrow highly ionized H-like emission possibly associated to the warm scattering gas on larger scales. In the following section, we attempt to quantify the location and physical properties of the soft X-ray emitting gas in Ark\,120. \subsection{Constraints from the He-like line triplets} The lack of intrinsic X-ray absorption towards Ark\,120 allowed a clean measurement of the emission from the He-like triplets. Given the constraints on the He-like line triplets, we can place an estimate on the density and subsequently infer the likely radial location of the emitting gas. The line ratios $G = (x + y + z) / w$ and $R = z / (x + y)$ give a measure of the temperature and density of the gas, where $z$ corresponds to the forbidden line, $(x + y)$ to the intercombination emission and $w$ to the resonance line (Porquet \& Dubau 2000). Taking the ratios of the fluxes of the line components measured from the O\,\textsc{vii} triplet (Table\,5), then $R=1.1\pm0.6$, resulting from a relatively equal contribution of the broad forbidden and intercombination components. This indicates that the gas is of relatively high density. From the calculations of Porquet \& Dubau (2000), this ratio corresponds to an electron density of $n_{\rm e}\sim 10^{11}$\,cm$^{-3}$. On the other hand the high G ratio, of $G=2.1\pm1.0$, indicates that the gas being photoionized rather than collisionally ionized, with a temperature of $T<10^6$\,K. Note that photoionization rather than photo-excitation appears the more dominant mechanism in Ark\,120, given the relative weakness of the higher order emission lines (see Fig 5, Kinkhabwala et al. 2002); for instance for O\,\textsc{vii} the ratio of the He$\alpha$ to He$\beta$ emission is $>10$. In contrast for the He-like triplet of Ne\,\textsc{ix}, the forbidden line appears to dominate over any intercombination emission (see Table 5), as generally the Ne\,\textsc{ix} triplet is more sensitive towards higher densities compared to O\,\textsc{vii}. Thus the R ratio of $R>3.3$ implies a limiting density of $n_{\rm e}\ls 2\times10^{11}$\,cm$^{-3}$. So for Ark\,120, the broad lined soft X-ray emitting gas appears consistent with a density in the range from $10^{11}\ls n_{\rm e}\gs 2\times10^{11}$\,cm$^{-3}$. The G ratio from the Ne\,\textsc{ix} triplet is $G>3.8$, consistent with a photoionized plasma. Note that the constraints on $R$ obtained from the N\,\textsc{vi} triplet are consistent with this, with a resulting lower limit on the density of $n_{\rm e}>10^{10}$\,cm$^{-3}$. \subsection{The Ionization State and Location of the Gas} Thus the above density would seem to imply an origin of the broad line emission consistent with the optical BLR (Davidson \& Netzer 1979). The ionization of the emitter can also be constrained, given the measured line flux ratio of O\,\textsc{vii} He$\alpha$\,/\,O\,\textsc{viii} Ly$\alpha$\,$\sim10$, e.g. Table~4. Indeed from fitting an \textsc{xstar} emission model with a density of $n_{\rm e}=10^{11}$\,cm$^{-3}$ to model the O\,\textsc{vii} emission (see Section\,6.4), an ionization parameter of $\logxi=0.5$ is required in order for the less ionized He-like emission to dominate over the H-like emission. From this an estimate of the radial distance of the emitter can be obtained via the definition of the ionization parameter, i.e. $r = (L_{\rm ion} / \xi n_{\rm e})^{1/2}$, where $L_{\rm ion}$ is the $1-1000$\,Rydberg luminosity and $n_{\rm e}$ is the electron density. From the broad-band UV to hard X-ray SED, the ionizing luminosity of Ark\,120 is $L_{\rm ion}\sim 10^{45}$\,erg\,s$^{-1}$ (Porquet et al. 2016). If we adopt a density of $n_{\rm e}=10^{11}$\,cm$^{-3}$ from the above considerations, then the emitting radius is $r=5\times10^{16}$\,cm (or $\sim0.01$\,pc), again consistent with typical BLR radii (Kaspi et al. 2005). Furthermore the distance to the optical BLR, inferred from time delays of the H$\beta$ line wrt to the continuum in Ark\,120 is $\tau\sim40$\,light days (Peterson \& Gaskell 1991, Wandel, Peterson \& Malkan 1999), equivalent to a distance of $10^{17}$\,cm. The radius of the emission can also be estimated from the O\,\textsc{vii} or Ne\,\textsc{ix} velocity widths of $\sigma\sim2000$\,km\,s$^{-1}$ or a FWHM of $\sim5000$\,km\,s$^{-1}$, adopting the more conservative lower value from the triplet deconstruction (see Table 5). Assuming a standard virial relation between the black hole mass and the radius $r$, of $3\sigma^{2}=GM/r$ and adopting a black hole mass of $1.5\pm0.2\times10^{8}$\Msun\ for Ark\,120 from reverberation mapping (Peterson et al. 2004), gives a radius of $r\sim10^{17}$\,cm, consistent with the above estimates. In comparison the optical H$\beta$ FWHM line width of Ark\,120 is $5850\pm480$\,km\,s$^{-1}$ (Wandel, Peterson \& Malkan 1999), which is similar to (or slightly smaller) than the typical X-ray broad line widths measured here. Furthermore the core of the 6.4\,keV iron K$\alpha$ emission line is also resolved in the simultaneous Chandra/HETG spectrum, with a FWHM width of $4500^{+2500}_{-1500}$\,km\,s$^{-1}$ (Nardini et al. 2016), consistent with the above estimates. Thus the observations suggest that the soft X-ray broad emission lines originating from a higher ionization phase of the AGN BLR, with radii in the typical range from $5\times10^{16}-10^{17}$\,cm. In contrast the H-like profiles measured in the spectra from N, O and Mg appear narrow and are unresolved by the RGS -- see Figure 7 and Table 5. The constraints on the density of the narrow lined emitting gas are not as tight as for the broad lines, with the Mg\,\textsc{xi} triplet measured from {\it Chandra} (forbidden dominating over intercombination with $R>3.2$) indicating $n_{\rm e}<10^{12}$\,cm$^{-3}$. Thus although we cannot directly measure the density of the narrow line emitting gas from line diagnostics, we can calculate its likely radial location from the limits on the velocity widths. If we take the line width for the resolved Mg\,\textsc{xi} forbidden line from the {\it Chandra} spectrum, with $\sigma=450^{+290}_{-180}$\,km\,s$^{-1}$, then the likely radius of the emitting gas is on pc scales. Note this is also consistent with the limits on the widths of the narrow lines seen in the RGS, e.g. for the O\,VIII Ly$\alpha$ line then $\sigma<470$\,km\,s$^{-1}$. For comparison, the expected distances of the torus and of the Narrow Line Region (NLR) are about 3\,pc and 100\,pc, using the following formula of Krolik \& Kriss (2001) and Mor, Netzer \& Elitzur (2009) respectively: \begin{equation} R_{\rm torus}\sim L_{\rm ion,44}^{1/2} ~~~~~~~~ (pc) \end{equation} \begin{equation} R_{\rm NLR} = 295 \times L_{46}^{0.47\pm 0.13} ~~~~ (pc). \end{equation} Thus, unlike for the broad line emitting gas, the narrow line X-ray emission appears consistent with radial locations commensurate with the pc scale torus or innermost NLR. This is also at a similar location to some of the soft X-ray warm absorbers inferred in Seyfert 1 AGN (see Tombesi et al. 2013 and references therein), which may imply that we are viewing similar ionized gas in Ark\,120 out of our direct line of sight. \subsection{The Covering Fraction and Geometry of the Gas} The luminosity of the soft X-ray line emission can also be used to calculate the global covering factor of the gas (see also Nucita et al. 2010). From the photoionization modeling, the normalization (or flux), $\kappa_{\rm xstar}$, of an emission component is defined by \textsc{xstar} (Kallman et al. 2004) in terms of: \begin{equation} \kappa_{\rm xstar} = f\frac{L_{38}}{D_{\rm kpc}^2} \end{equation} where $L_{38}$ is the ionizing luminosity in units of $10^{38}$\,erg\,s$^{-1}$, $D_{\rm kpc}$ is the distance to the AGN in kpc. Here $f$ is the covering fraction of the gas with respect to the total solid angle, where $f = \Omega / 4\pi$. For a spherical shell of gas, $f=1$, while $L$ is the quasar ionizing luminosity that illuminates the photoionized shell. The ionizing luminosity obtained from the best fit spectral model and corrected for the intervening Galactic absorption, is $L=4.5\times10^{44}$\,erg\,s$^{-1}$ over the 1--1000\,Rydberg band\footnote{Note that the ionizing luminosity returned from a broad-band fit to Ark\,120 with the \xmm\ EPIC and OM data, obtained from the \textsc{optxagn} disk + corona model (Porquet et al. 2016, in preparation), is $\sim 10^{45}$\,erg\,s$^{-1}$.}. Thus for Ark\,120 with $D=133$\,Mpc and for a spherical shell, the expected {\sc xstar} normalization from above is $\kappa_{\rm xstar}=2.5\times10^{-4}$. Hence for a given column density of gas, this sets the total luminosity of the soft X-ray photoionized emission. As a first step, model emission line spectra were then generated with this overall normalization factor within \textsc{xstar}, which gives the predicted emission originating from a fully covering spherical shell of gas (with $f=1$), illuminated by an AGN ionizing luminosity of $L$. The ionization of the gas was fixed at $\logxi=0.5$, consistent with the \textsc{xstar} fits to the lines (see below). Taking as an example the case of the strong O\,\textsc{vii} He$\alpha$ broad emission line observed in the RGS spectrum, we then compared the observed line luminosity to that predicted by the \textsc{xstar} model and then used the ratio of the observed to predicted line luminosity to calculate the global covering fraction of the gas. From the Ark\,120 spectrum, the observed luminosity of the broad O\,\textsc{vii} emission is $L_{\rm OVII}=4.0\pm1.2\times10^{41}$\,erg\,s$^{-1}$. In comparison, for a column density of $N_{\rm H}=10^{21}$\,cm$^{-2}$, the O\,\textsc{vii} luminosity predicted by the \textsc{xstar} model for a fully covering shell of gas is $1.2\times10^{42}$\,erg\,s$^{-1}$. The ratio of the observed to predicted luminosity then gives the geometric covering fraction of the emitter, which for $N_{\rm H}=10^{21}$\,cm$^{-2}$ is $f=0.33\pm0.10$. Similarly, for a higher column of $N_{\rm H}=10^{22}$\,cm$^{-2}$, the predicted O\,\textsc{vii} luminosity of a spherical shell is higher, with $3.6\times10^{42}$\,erg\,s$^{-1}$ and thus the covering fraction is then lower, with $f=0.11\pm0.03$. To provide a more quantitative estimate of the covering fraction, the RGS spectrum was then fitted with the \textsc{xstar} emission models. In order to reproduce the strongest emission lines present in the Ark\,120 spectrum, three different photoionized emission zones were required and their properties are summarized in Table\,6. Two of these emission zones appear to be broadened (with $\sigma=3000$\,km\,s$^{-1}$); these are required to model (i) the broad O\,\textsc{vii} emission (as well as at N\,\textsc{vi}) and (ii) the broad Ne\,\textsc{ix} emission, with the lower ionization zone (with $\log \xi=0.5\pm0.1$) responsible for O\,\textsc{vii}. A third, highest ionization zone ($\log \xi=2.3\pm0.4$) with a narrow velocity width ($\sigma=300$\,km\,s$^{-1}$) is responsible for the narrow H-like lines, such as from O\,\textsc{viii} Ly-$\alpha$. The column density of the gas in emission is not known a priori (due to the lack of absorbing gas) and is highly degenerate with the emission normalization. Thus instead of directly fitting both the column and the normalization of the \textsc{xstar} emission zones, the column density was varied over the range $3\times10^{20}<N_{\rm H}<10^{22}$\,cm$^{-2}$, adopting 8 different values and at each fixed $N_{\rm H}$ value the spectrum was refitted to obtain the normalization of the \textsc{xstar} component. Thus by comparing the observed normalization ($\kappa_{\rm obs}$) of a fitted emission component with the predicted value from equation\,3 ($\kappa_{\rm xstar}$), the covering fraction for a given column can be calculated by $f=\kappa_{\rm obs}/\kappa_{\rm xstar}$. The resulting plot of covering fraction vs. column for the O\,\textsc{vii} emission zone is shown in Figure\,10, noting that similar results are also found for the Ne\,\textsc{ix} zone and the (narrow) O\,\textsc{viii} zone (see Table 6). The overall trend is for the covering fraction of the gas to decrease with increasing column. This would be expected as increasing the column density of the gas clouds increases their soft X-ray luminosity, requiring the overall covering fraction to decrease in order to compensate. A minimum column density of $N_{\rm H}>3\times10^{20}$\,cm$^{-2}$ is required to reproduce the required O\,\textsc{vii} luminosity, if the gas is fully covering the AGN with $f=1$. This fully covering scenario appears less likely, as the upper-limit to the line of sight column of soft X-ray warm absorbing gas towards Ark\,120 is at least factor of $\times10$ lower; with $N_{\rm H}<3\times10^{19}$\,cm$^{-2}$ for $\logxi=1$, see Section 5.5. Instead the covering fraction is likely to be lower, with no significant distribution of gas along the line of sight, allowing the column density out of the direct line of sight to be higher. Indeed once the column density approaches $N_{\rm H}=10^{22}$\,cm$^{-2}$, then the covering reaches a limiting value of $f=0.1$. This corresponds to a likely {\it minimum} covering fraction of gas as increasing the column density above $N_{\rm H}>10^{22}$\,cm$^{-2}$ has little effect on the total line luminosity; i.e at these columns and higher, the emitting clouds become optically thick at soft X-rays, with little change in the resulting soft X-ray line luminosity. \subsection{X-ray Broad and Narrow Line Regions} Thus the above calculations, in order to reproduce the broad soft X-ray line emission, requires the emitting gas to have a typical column of up to $N_{\rm H}\sim10^{22}$\,cm$^{-2}$, with a covering fraction of at least 10\% of $4\pi$ steradian and for the gas to lie out of the direct line of sight. The mass of the emitting gas implied from these calculations is $\sim 0.1$\,M$_{\odot}$, for $N_{\rm H}=10^{22}$\,cm$^{-2}$ and $f=0.1$. Overall this scenario is similar to what is usually inferred from studies of the optical--UV BLR (Baldwin et al. 1995, Gaskell 2009), with the distribution of emitting BLR clouds likely non-spherical and following a flattened or disk-like geometry (Krolik et al. 1991, Eracleous \& Halpern 2003). The latter distribution of emitting clouds can also account for the lack of absorption towards Ark\,120 if the gas lies out of the direct line of sight and our view of the AGN is relatively pole-on compared to the plane of the disk. This may also be consistent with the morphology of the host spiral galaxy being relatively pole-on in Ark\,120 (Nordgren et al. 1995). The opposite may be the case in AGN where variable X-ray obscuration events occur, when the line of sight can be intercepted by compact clouds consistent with observed BLR distances. Such X-ray obscuration events have been observed several Seyfert galaxies, such as in Mrk 335 (Longinotti et al. 2013), NGC 3227 (Lamer et al. 2003), NGC 1365 (Risaliti et al. 2005b, Braito et al. 2014), NGC 4051 (Terashima et al. 2009, Lobban et al. 2011), NGC\,3516 (Turner et al. 2008) and NGC 5548 (Kaastra et al. 2014). In contrast to the broadened emission, the kinematics of the narrow soft X-ray emission lines are consistent with arising from pc scale distances or higher. The typical column density and covering fraction of the narrow emission line gas is similar to broad lined gas; for a column of $N_{\rm H}=1\times10^{21}$\,cm$^{-2}$, the covering fraction is consistent with $f=1$, while the covering fraction decreases to $\sim 10$\% for $N_{\rm H}=1\times10^{22}$\,cm$^{-2}$. Thus the derived columns and covering fractions are thus likely too low to be commensurate with emission from a Compton thick pc-scale torus. However the origin of the gas in emission may arise from the AGN Narrow Lined Region (NLR) and is in general agreement with the typical columns and spatial locations of matter inferred along the line of sight in the warm absorbers that are observed in many other Seyfert 1 AGN (e.g. Mrk 509, Kaastra et al. 2012). Indeed the kinematics of the narrow emission lines, in terms of the velocity widths and a tentative indication of modest blueshift (both of the order of a few hundred km\,s$^{-1}$), may imply we are observing emission from a pc scale outflow, but viewed out of the line of sight towards Ark\,120.

1607.01062

1607

1607.08005_arXiv.txt

We conducted a deep narrow-band imaging survey with the Subaru Prime Focus Camera on the Subaru Telescope and constructed a sample of Ly$\alpha$ emitters (LAEs) at $z=2.53$ in the UDS-CANDELS field where a sample of H$\alpha$ emitters (HAEs) at the same redshift is already obtained from our previous narrow-band observation at NIR. The deep narrow-band and multi broadband data allow us to find LAEs of stellar masses and star-formation rates (SFRs) down to $\gtrsim$$10^8$ M$_\odot$ and $\gtrsim$0.2 M$_\odot$/yr, respectively. We show that the LAEs are located along the same mass-SFR sequence traced by normal star-forming galaxies such as HAEs, but towards a significantly lower mass regime. Likewise, LAEs seem to share the same mass--size relation with typical star-forming galaxies, except for the massive LAEs, which tend to show significantly compact sizes. We identify a vigorous mass growth in the central part of LAEs: the stellar mass density in the central region of LAEs increases as their total galaxy mass grows. On the other hand, we see no Ly$\alpha$ line in emission for most of the HAEs. Rather, we find that the Ly$\alpha$ feature is either absent or in absorption (Ly$\alpha$ absorbers; LAAs), and its absorption strength may increase with reddening of the UV continuum slope. We demonstrate that a deep Ly$\alpha$ narrow-band imaging like this study is able to search for not only LAEs but also LAAs in a certain redshift slice. This work suggests that LAEs trace normal star-forming galaxies in the low-mass regime, while they remain as a unique population because the majority of HAEs are not LAEs.

The redshift interval of $z$=2.1--2.6 is the key epoch for us to understand the formation mechanisms of star-forming galaxies. In this era, the cosmic star-formation density is peaked \citep{Hopkins:2006}, galaxy morphology is still under construction \citep{Papovich:2005}, and gaseous flows (feeding and feedback) are more prevalent compared to lower redshift counterparts \citep{Yabe:2015}. Galaxies at this epoch allow us to study the physical origins of these processes based on their strong emission lines such as \ha$\lambda6563$, \oiii$\lambda5007$, \oii$\lambda3727$, and \lya$\lambda1216$ with the optical and near-infrared imagers and spectrographs. For example, we can directly investigate the physical properties of \lya\ emitters (LAEs) in this redshift range such as gaseous metallicities and ionization parameters \citep{Nakajima:2013}. We can also compare the characteristics of LAEs with those of \ha\ emitters (HAEs) \citep{Hayes:2010,Oteo:2015} up to $z\sim2.5$ which are relatively unbiased sample of star-forming galaxies. At higher redshifts, a comparison can be made with Lyman break galaxies (LBGs; \citealt{Verhamme:2008,Malhotra:2012}). Such comparison analyses should unveil the physical characteristics of LAEs. In addition, the observations of nearby LAE analogues also help us resolve the physical mechanism of \lya\ radiative transfer in detail \citep{Kunth:1998,Atek:2009b,Hayes:2013,Hayes:2015,Rivera:2015}. The properties of LAEs provide information on the reionisation history in the early Universe (e.g. \citealt{Dawson:2007,Kashikawa:2006,Stark:2010}), although the intrinsic properties of LAEs may also change with redshift \citep{Nilsson:2011}. The past analyses find that LAEs tend to be younger, less massive, and less dusty star-forming galaxies (e.g. \citealt{Gawiser:2006,Nilsson:2007,Ouchi:2008,Guaita:2011}). The escape fraction of \lya\ photons and the ionization parameter are both high in these systems, compared to other star-forming galaxies \citep{Nakajima:2012,Nakajima:2013,Oteo:2015}. Note that a minor fraction of LAEs can constitute massive and red galaxies (see also \citealt{Oteo:2012, Bridge:2013,Sandberg:2015,Finkelstein:2015,Taniguchi:2015}). Furthermore, LAEs tend to lie above the tight star-formation rate (SFR) vs. stellar mass relation called the main sequence of star-forming galaxies (\citealt{Vargas:2014, Hagen:2014}, but see also \citealt{Nilsson:2011b,Song:2014,Kusakabe:2015,Finkelstein:2015}). LAEs tend to have more compact star-forming regions seen in the rest-frame UV light compared to LBGs, and their angular sizes of UV continuum are independent of redshift at $z<6$ \citep{Malhotra:2012}. A more systematic study has been carried out by \citet{Hathi:2016}, who compare the physical properties among non-LAEs, LAEs with small equivalent width (EW$<20$ \AA), and those with EW$>20$ \AA. They have shown that the UV slope and the UV magnitude at 1500 \AA, M$_{1500}$, of star-forming galaxies with and without \lya\ emission lines are similar for a given i-band magnitude limit ($<25$ mag). In addition, \citet{Hagen:2016} have found no statistical difference in physical and morphological parameters such as specific SFR, SFR surface density, half-light radius and \oiii\ equivalent width between LAEs and non-LAEs at $z\sim2$, based on the sample limited to the rest-frame optical line flux of $\gtrsim10^{-17}$ erg/s/cm$^2$. However, a careful consideration is needed to discuss representative characteristics of LAEs. The excess of SFRs in LAEs with respect to the main sequence can be interpreted in such a way that a large amount of ionizing photons from young and active starbursts may allow \lya\ photons to escape from the systems more easily. On the other hand, various previous studies suggest that the properties of LAEs are quite diverse, which may imply that the escape of \lya\ photons is a stochastic process rather than an ordered duty cycle (see also \citealt{Nagamine:2010}). Moreover, the escape of \lya\ photons depends not only on the energy input from star-forming regions or active galactic nuclei but also on the amount of dust and circumgalactic medium \citep{Verhamme:2012,Yajima:2012,Shibuya:2014a,Shibuya:2014b}. Larger gas covering fraction leads to higher line of sight extinctions. \citet{Reddy:2016b} have found that those quantities correlates with UV continuum slope in the sense of redder UV continua tend to have higher covering fraction of \hi\ and dust. In addition, the effect of stellar absorption is also a non-negligible factor \citep{Pena:2013}. Another concern is that all the previous LAE studies have been limited by the \lya\ limiting flux that can be reached in the observations. The \lya\ luminosity depends on galaxy populations, and in fact most of the very bright LAEs ($\gtrsim3\times10^{43}$ erg/s) host active galactic nuclei \citep{Ouchi:2008,Konno:2016,Sobral:2016}. For example, we may find LAEs on the main sequence of normal star-forming galaxies when we conduct deeper observations. If most of the LAEs have normal specific SFRs similar to those of typical star-forming galaxies, the escape of the \lya\ photons cannot be due simply to their high star-formation activity, and some other factors would be more dominant. To understand the nature of LAEs, the dual emitter survey of HAEs and LAEs at the same redshift is very powerful. \citet{Hayes:2010} first reported that only six out of 55 \ha-selected galaxies were detected by a narrowband filter for \lya\ line at $z=2.2$. They also found very low \lya\ photon escape fractions of $\sim5\%$, which is consistent with another recent dual survey of HAEs and LAEs \citep{Matthee:2016,Sobral:2016}. Those studies noted that it is difficult to account for such a low \lya\ escape fraction only by dust attenuation (see also \lya\ line observations of nearby galaxies e.g. \citealt{Atek:2008}). All these studies point towards the fact that most of the star-forming galaxies do not actually show \lya\ emission lines that are strong enough to be detected by a standard narrowband imaging observation. What if we go significantly deeper in \lya\ narrowband imaging observation? The fraction of dual emitters (HAEs and LAEs) should depend on the survey depths. Indeed, the most of the past \lya\ line imaging surveys with NB$_\mathrm{5\sigma}\lesssim25$ mag do not reach continua at \lya\ line wavelength for many \ha-selected galaxies at $z>2$, in particular for dusty galaxies. Deeper \lya\ observation will give us more comprehensive insights into the \lya\ emissivities of high redshift galaxies. Our project named MAHALO-Subaru (Mapping HAlpha Line of Oxygen with Subaru) has successfully mapped out \ha, \oii, and \oiii\ emitters at $z$=0.4--3.8 in 14 overdense regions (e.g. \citealt{Kodama:2004, Koyama:2013a,Hayashi:2014}) and some blank fields in COSMOS- and UDS- CANDELS fields (e.g. \citealt{Tadaki:2013, Suzuki:2015}). Some follow-up spectroscopic surveys have been conducted by \citet{Shimakawa:2014,Shimakawa:2015,Shimakawa:2015b}. We have now recently conducted a deep narrow-band survey of \lya\ emitters with Suprime-Cam in coordination with the existing companion survey of \ha\ emitters with MOIRCS \citep{Tadaki:2013}. This enables us to probe star-forming galaxies over a wider mass range and across various environments at $2<z<2.6$, where both \ha\ and \lya\ lines can be accessed with the ground-based telescopes. This {\it Paper} studies the LAEs at $z=2.53$ in the UDS-CANDELS field based on the deep NB imaging observations with Suprime-Cam, as the first such targets. We mainly focus on star-forming activities and stellar sizes of the LAEs, as compared to HAEs. The details of the observations and the analyses can be found in \S2. \S3 presents the star-forming activities and stellar sizes of LAEs, and compares them with those of HAEs, which are more representative, massive star-forming galaxies at the same redshift. Discussion is given in \S4, and we conclude this work in \S5. We assume the cosmological parameters of $\Omega_M$=0.3, $\Omega_\Lambda$=0.7 and $H_0$=70 km/s/Mpc and employ a \citet{Chabrier:2003} stellar initial mass function (IMF). The AB magnitude system \citep{Oke:1983} is used throughout this paper. \newpage

We have conducted an optical narrow-band imaging survey with the Suprime-Cam to search for LAEs at $z=2.5$ in the UDS-CANDELS field where we have already identified many HAEs in the same redshift slice by our previous narrow-band imaging survey at NIR. Together with existing WFC3/IR data from the HST and multi-band photometries from various large programs, we identify 50 LAE candidates at $z$=2.53 down to stellar mass of 10$^8$ \msun and SFR of 0.2 \msun/yr. Our sample are limited by \lya\ luminosities ($>4.40\times10^{41}$ erg/s in 1.5 arcsec aperture diameter), \lya\ equivalent widths ($>15$ \AA), and H$_\mathrm{F160W}$ band magnitude ($<26.89$ mag). We combine 37 HAEs from our previous work \citep{Tadaki:2013}, and investigate the physical properties of LAEs and HAEs such as their SFRs and sizes. We compare star-forming activities and stellar angular sizes of LAEs with those of HAEs and other star-forming galaxies in the literature. The F160W/IR images for galaxies at $z=2.53$ are not affected by \hb\ or \oiii\ lines, which thus provide us with the accurate stellar masses and the sizes of the LAEs and HAEs. The depth of our narrow-band and B-band data allows us to explore star-forming galaxies with lower stellar masses down to $\gtrsim$10$^8$ \msun. As a result, we find that LAEs follow the extrapolated line of the mass--SFR relation (star-forming main sequence) to the lower-mass end. The result is in good agreement with and complementary to \citet{Hagen:2016} which have reported no difference between LAEs and controlled sample for comparison on the mass--SFR diagram at the similar redshift. The \lya\ luminosities and EWs do not strongly correlate with sSFR, indicating that star-forming activities do not contribute directly to the \lya\ photon escape mechanism. As suggested by recent studies (e.g. \citealt{Yajima:2012,Shibuya:2014b,Reddy:2016b}), gas covering fraction and dust extinction would be the primary key factors. Indeed, we identify \lya\ absorption features in a part of the HAEs individually, which tend to be dustier and more massive systems. Since NB428 filter can cover \lya\ emissions of the existing HAEs discovered by the past narrow-band (NB2315) imaging at NIR, the combination of these two narrow-band filter allow us to investigate \lya\ emissivities of HAEs. However, we detect significant flux excesses in only two HAEs. Six out of 37 HAEs have positive \lya\ EWs, of which 2 show significant \lya\ flux excesses at more than 3 sigma levels. Such a small fraction of LAEs among HAEs is consistent with the past similar studies \citep{Hayes:2010,Matthee:2016}, although it should be noted that the percentage should depend on the survey depths, photometric aperture size, physical properties of the parent HAE samples, and so on \citep{Matthee:2016}. For example, photometric aperture size of 1.5 arcsec diameter used in this work is a half of that used in \citet{Matthee:2016}. We thus focus more on the \lya\ emissivity or its escape fraction along the line of sight. Instead, we find flux deficits in the narrow-band for a larger fraction of our HAE sample. The flux-limited LAEs and HAEs are quite distinctively distributed on the EW$_\mathrm{Ly\alpha}$ versus stellar mass or EW$_\mathrm{Ly\alpha}$ versus E(B$-$V) diagrams. We conclude that LAEs have similar general properties as those of normal star-forming galaxies. However, they also remain as a unique population (low-mass young galaxies) because the majority of HAEs are not LAEs. LAEs (low-mass star-forming galaxies) seem to share the same mass-size relation of massive star-forming galaxies within errors. On the other hand, four out of five massive LAEs at log(M$_\star$/\msun)$>$9.5 have compact structures and are located close to the mass-size relation of early-type galaxies rather than that of late-type galaxies reported by \citet{Wel:2014}. Moreover, they have low sSFRs and high mass surface densities. Such massive LAEs have experienced nuclear starbursts in the past, which clear the surrounding gas by a galactic wind, allowing the \lya\ photon to escape from the systems as detected in our narrow-band \lya\ imaging. Finally, we demonstrate the unique narrow-band technique to search for \lya\ absorbers (LAAs), which are observed as showing flux deficits in the narrow-band instead of flux excesses. Whilst this method can effectively trace LAAs in a certain redshift slice, we should note that the sample is limited to a narrow range of EW$_\mathrm{Ly\alpha}$ in absorption due to observational limitations. In order to generalise our results to normal star-forming galaxies, we will need to investigate other types of galaxy populations than LAEs at the same redshift and in the same stellar-mass range, and compare their physical properties with those of LAEs. We are now conducting systematic ultra-deep narrow-band imaging survey of low-mass, low-luminosity HAEs (Kodama et al.) for this purpose.

1607.08005

1607

1607.06829_arXiv.txt

In the presence of a strong magnetic field a stellar equilibrium configuration, aided by the Lorentz force, can support a larger mass than a non-magnetic one. This has been considered a possible explanation of the super-Chandrasekhar mass white dwarfs giving rise to over-luminous Type-Ia supernovae. We present here linear and non-linear perturbation studies of such strongly magetised configurations and show that axisymmetric configurations with poloidal or toroidal fields are unstable. The numerical evolution of the perturbations shows instability after about an Alfv\'en crossing time. This time scale is very short for the magnetically supported super-Chandrasekhar mass white dwarfs. Uniform rotation about the symmetry axis can reduce the growth rate but can not stabilize the super-massive configurations. It is concluded that long-lived super-Chandrasekhar mass white dwarfs supported by magnetic field are unlikely to occur in Nature.

White dwarfs, compact object supported by electron degeneracy pressure, have the well known maximum mass limit of 1.4 M$_\odot$ \citep{Chandrasekhar1931}. If accretion raises the mass above this Chandrasekhar limit, the white dwarf is no longer able to support itself against gravitational collapse, and the resulting rapid contraction leads to a Type-Ia supernova. The characteristic mass limit sets the standard properties of the Type-Ia supernova. Recently however a few cases have been observed of over-luminous Type Ia supernovae that require white dwarfs well above this mass limit ($\mbox{$\stackrel{>}{_{\sim}}$}$ 2 M$_\odot$) to explain their properties \citep{howel06,hicken_07,yamanaka_09, Scalzo+2010, Tanaka+2010, Silverman+2011, Taubenberger+2011}. \cite{das_m12} explored the possibility of super-Chandrasekhar mass configurations arising from quantum mechanical modification of the degenerate electron equation of state in the presence of ultra strong internal magnetic fields. The effect of this turns out to be sub-dominant to that of the Lorentz force; the latter by itself can raise the maximum mass well above the Chandrasekhar limit \citep{Ostiker_Hartwick68, Bera+Bhattacharya2014}. The maximum mass of the magnetically supported configurations is dependent on the field geometry. Among axisymmetric structures, the maximum mass is about 1.9 M$_\odot$ for pure poloidal field \citep{Bera+Bhattacharya2014, das_m15, Franzon+Schramm2015} and more than 5 M$_\odot$ for pure toroidal field \cite{Bera+Bhattacharya2016}, with intermediate values for mixed field configurations. These limits refer to equilibrium structures without consideration of stability. In this paper, we study the stability of these equilibrium configurations. Magnetic field is ubiquitous in the compact stars, be it neutron stars or white dwarfs. The highest magnetic field measured at the neutron star surface is $\sim 10^{15}$ G and that of a white dwarf is $\sim 10^9$ G \citep{Schmidt+2003}. While the field external to the star is primarily poloidal, \cite{Prendergast1956} suggests that in the interior both poloidal and toroidal fields must be present to ensure long term stability. The consideration of the minimum energy principle \citep{Bernstein+1958} indicates that equilibrium configurations with pure poloidal \citep{marke73} or pure toroidal \citep{tayler1973} magnetic field are unstable. For such field geometries, perturbations in the matter and the magnetic field close to the neutral line (viz. the locus of vanishing magnetic field in the stellar interior, enclosed by field lines) can generate states of lower energy than the unperturbed configuration. Perturbation of the system would therefore drive it to a new configuration by rearranging the magnetic field and matter. This magnetic instability is intrinsic to the field geometry, even if the magnetic energy is small compared to the thermal and the gravitational energy of the configuration (see e.g. \cite{Flowers+Ruderman77}). \cite{tayler1973} and \cite{Acheson1978} show that for a pure toroidal configuration the non-axisymmetric azimuthal mode $m = 1$ is the dominant instability mode with very short instability time scale (Alfv\'en crossing time). The non-linear evolution of the magnetic configurations with pure poloidal and pure toroidal field shows instability with growth time comparable to the Alfv\'en time of the configuration \citep{Braithwaite+Spruit2006, Braithwaite2006b, Bonanno+Uprin2013a, Bonanno+Uprin2013b, Bonanno+Uprin2013c, IbanezMejia+Braithwaite2015}. Configurations that show long term dynamical stability in numerical experiments with a stably stratified star contain comparable amounts of energy in poloidal and toroidal components \citep{Braithwaite+Nordlund2006, braithwaite09}. Numerically evolved axisymmetric or non-axisymmetric stable structures have been found for configurations with non-barotropic (stably stratified) equation of state \citep{mitchell+2015} and a helical initial field distribution with random or mixed poloidal-toroidal field \citep{Braithwaite2008}. Numerical studies of the evolution of neutron stars with pure poloidal fields in general relativistic formalism have been carried out by \cite{Lasky+2011, Ciolfi+2011, Ciolfi+Rezzolla2012}. These studies also indicate the presence of instability near neutral line. Magnetars, a set of neutron stars with a strong surface magnetic field ($\sim 10^{15}$ G), show repeated gamma-ray flares. Quasi-periodic oscillations (QPOs) observed in the tail of giant flares provide evidence of neutron star oscillations \citep{Isreal+2005, Strohmayer+Watts2005}. Such oscillations of strongly magnetized neutron stars have been investigated by various authors for both axisymmetric \citep{Glampedakis+2006, Lee2008, gabler+13, gabler+12} and non-axisymmetric \citep{Lander+2010, Lander+Jones2011a, Lander+Jones2011b, Asai+2015, Asai+2016} modes. In this paper, we study the linear and non-linear evolution of perturbed variables of a magnetized equilibrium structure. The perturbation equations and the methods of evolution are described in Section~\ref{eq+method}. In Section~\ref{results} we present in brief the results obtained, which we discuss in Section~\ref{discussion}. Our conclusions are summarized in Section~\ref{conclusion}.

\label{conclusion} In this paper, we have studied the linear and non-linear evolution of perturbations to an axisymmetric, strongly magnetized object. Our main results are: \begin{enumerate} \item Axisymmetric magnetic configurations with pure poloidal or pure toroidal field suffer from magnetic instabilities with time scale comparable to the Alfv\'en crossing time ($\tau_A$) of the configuration. As the Alfv\'en crossing time is inversely proportional to the average magnetic field of the configuration, structures with ultra strong magnetic fields, with very short Alfv\'en crossing times, are strongly unstable. \item In the case of rotating magnetic white dwarfs, the instability growth rate reduces as the rotation speed increases. However this is insufficient to fully stabilize magnetically supported configurations near their mass limits. \item Magnetically supported super-Chandrasekhar mass white dwarfs require extremely strong magnetic fields in the interior and are hence susceptible to instabilities with a very short growth time scale (typically less than a second). Instabilities of this nature may in fact prevent the formation of such objects. \end{enumerate}

1607.06829

1607

1607.02707_arXiv.txt

{Infrared-faint radio sources~(IFRS) are a class of radio-loud~(RL) active galactic nuclei~(AGN) at high redshifts~($z\geq 1.7$) that are characterised by their relative infrared faintness, resulting in enormous radio-to-infrared flux density ratios of up to several thousand. } {Because of their optical and infrared faintness, it is very challenging to study IFRS at these wavelengths. However, IFRS are relatively bright in the radio regime with 1.4\,GHz flux densities of a few to a few tens of mJy. Therefore, the radio regime is the most promising wavelength regime in which to constrain their nature. We aim to test the hypothesis that IFRS are young AGN, particularly GHz peaked-spectrum~(GPS) and compact steep-spectrum~(CSS) sources that have a low frequency turnover.} {We use the rich radio data set available for the Australia Telescope Large Area Survey fields, covering the frequency range between 150\,MHz and 34\,GHz with up to 19~wavebands from different telescopes, and build radio spectral energy distributions~(SEDs) for 34~IFRS. We then study the radio properties of this class of object with respect to turnover, spectral index, and behaviour towards higher frequencies. We also present the highest-frequency radio observations of an IFRS, observed with the Plateau de Bure Interferometer at 105\,GHz, and model the multi-wavelength and radio-far-infrared SED of this source.} {We find IFRS usually follow single power laws down to observed frequencies of around 150\,MHz. Mostly, the radio SEDs are steep ($\alpha < -0.8$; $74^{+6}_{-9}$\%), but we also find ultra-steep SEDs ($\alpha < -1.3$; $6^{+7}_{-2}$\%). In particular, IFRS show statistically significantly steeper radio SEDs than the broader RL AGN population. Our analysis reveals that the fractions of GPS and CSS sources in the population of IFRS are consistent with the fractions in the broader RL AGN population. We find that at least $18^{+8}_{-5}$\% of IFRS contain young AGN, although the fraction might be significantly higher as suggested by the steep SEDs and the compact morphology of IFRS. The detailed multi-wavelength SED modelling of one IFRS shows that it is different from ordinary AGN, although it is consistent with a composite starburst-AGN model with a star formation rate of 170\,$M_\odot$\,yr$^{-1}$.} {}

\label{introduction} Infrared-faint radio sources~(IFRS) are comparatively bright radio sources with a faint or absent near-infrared counterpart. They were serendipitously discovered in the Chandra Deep Field-South~(CDFS) by \citet{Norris2006} in the Australia Telescope Large Area Survey~(ATLAS) 1.4\,GHz map and the co-located \textit{Spitzer} Wide-area Infrared Extragalactic Survey~(SWIRE; \citealp{Lonsdale2003}) infrared~(IR) map. Based on the SEDs of ordinary galaxies, it was expected that every object in the deep radio survey (rms of 36\,$\mu$Jy\,beam$^{-1}$ at 1.4\,GHz in CDFS) would have a counterpart in the SWIRE survey~(rms of $\sim 1\,\mu$Jy at 3.6\,$\mu$m). However, \citeauthor{Norris2006} found 22~radio sources with 1.4\,GHz flux densities of a few or a few tens of mJy without 3.6\,$\mu$m counterpart and labelled them as IFRS. Later, IFRS were also found in the European Large Area IR space observatory Survey South~1~(ELAIS-S1) field, the \textit{Spitzer} extragalactic First Look Survey~(xFLS) field, the Cosmological Evolution Survey~(COSMOS) field, the European Large Area IR space observatory Survey North~1~(ELAIS-N1) field, and the Lockman Hole field~\citep{Middelberg2008ELAIS-S1,GarnAlexander2008,Zinn2011,Banfield2011,Maini2013submitted}, resulting in around 100~IFRS known in deep fields.\par While IFRS were originally defined as radio sources without IR counterpart in the first works, \citet{Zinn2011} set two criteria for the survey-independent selection of IFRS: \begin{enumerate}[(i)] \item radio-to-IR flux density ratio $S_{1.4\,\textrm{GHz}}/S_{3.6\,\mu\textrm{m}} > 500$~, and \item 3.6\,$\mu$m flux density $S_{3.6\,\mu\textrm{m}} < 30\,\mu$Jy~. \end{enumerate} The first criterion accounts for the enormous radio-to-IR flux density ratios resulting from the solid radio detection and the IR faintness. These ratios identify IFRS as clear outliers. The second criterion selects objects at cosmologically significant redshifts because of cosmic dimming or heavily obscured objects.\par \citet{Collier2014} followed a different approach than used in the previous studies and searched for IFRS based on shallower data, but in a much larger area. Using the Unified Radio Catalog~(URC; \citealp{Kimball2008}) based on the NRAO VLA Sky Survey~(NVSS; \citealp{Condon1998}) and IR data from the all-sky Wide-Field Infrared Survey Explorer~(WISE; \citealp{Wright2010}), they found 1317~IFRS fulfilling both selection criteria from \citet{Zinn2011}. Whereas some of the IFRS in deep fields are lacking a 3.6\,$\mu$m counterpart, all IFRS from the catalogue compiled by \citeauthor{Collier2014} have a detected 3.4\,$\mu$m counterpart. Also, these sources are on average radio-brighter than the IFRS in deep fields.\par Since the first IFRS were identified, it has been argued that these objects might be radio-loud~(RL) active galactic nuclei~(AGN) at high redshifts~($z\gtrsim 1$), potentially heavily obscured by dust~\citep{Norris2006,Norris2011}. Whereas other explanations like pulsars have been ruled out~\citep{Cameron2011}, the suggested high redshifts of IFRS have been confirmed by \citet{Collier2014} and \citet{Herzog2014}; all spectroscopic redshifts are in the range $1.7\lesssim z \lesssim 3.0$. The first two very long baseline interferometry~(VLBI) detections of IFRS were carried out by \citet{Norris2007} and \citet{Middelberg2008IFRS_VLBI} who targeted six IFRS in total and show that at least some IFRS have high brightness temperatures, indicating the presence of an AGN. Recently, \citet{Herzog2015a} found compact cores in the majority of IFRS based on a large sample of 57~sources. \citet{Middelberg2011} show that IFRS have significantly steeper radio SEDs (median index\footnote{The spectral index~$\alpha$ is defined as $S\propto \nu^\alpha$ throughout this paper where $S$ is the flux density and $\nu$ the frequency.} of $-1.4$ between 1.4\,GHz and 2.4\,GHz) than ordinary AGN.\par An overlap between the populations of IFRS on the one hand and GHz peaked-spectrum~(GPS) and compact steep-spectrum~(CSS) sources on the other hand is suggested and found by \citet{Middelberg2011}, \citet{Collier2014} and \citet{Herzog2015a}. GPS sources are very compact and powerful AGN with linear sizes below 1\,kpc, showing a turnover in their radio spectral energy distribution~(SED) at frequencies of 500\,MHz or higher. CSS sources are similarly powerful, but are more extended (linear sizes of a few or a few tens of kpc) and show their turnover at frequencies below 500\,MHz~(e.g.\ \citealp{ODea1998,Randall2011}). Further, CSS sources are characterised by their steep radio SEDs ($\alpha \lesssim -0.5$). GPS and CSS sources are usually considered to be young versions of extended radio galaxies, but it has also been suggested that they are frustrated AGN confined by dense gas~\citep{ODea1991} or dying radio sources~\citep{Fanti2009b}.\par Modelling the multi-wavelength SED of IFRS was accomplished by \citet{GarnAlexander2008}, \citet{Huynh2010}, \citet{Herzog2014}, and \citet{Herzog2015b}, and shows that these sources can only be modelled as high-redshift RL AGN, potentially suffering from heavy dust extinction. The strong link between IFRS and high-redshift radio galaxies~(HzRGs)---first suggested by \citeauthor{Huynh2010} and \citet{Middelberg2011} and later emphasised by \citet{Norris2011}---has also been found in the modelling by \citet{Herzog2015b}. HzRGs are massive galaxies at high redshifts ($1\leq z \leq 5.2$) which are expected to be the progenitors of the most massive elliptical galaxies in the local universe (e.g.\ \citealp{Seymour2007,deBreuck2010}). They host AGN and undergo heavy star forming activity. IFRS have a significantly higher sky density than HzRGs (a few IFRS per square degree versus around 100~HzRGs known on the entire sky) and are suggested to be higher-redshift or less luminuous siblings of these massive galaxies.\par The correlation between $K$~band magnitude and redshift has been known for radio galaxies~(e.g.\ \citealp{Lilly1984,Willott2003,Rocca-Volmerange2004}) for three decades and was used to find radio galaxies at high redshifts. In particular, HzRGs were also found to follow this correlation~\citep{Seymour2007}. Although IFRS are on average fainter than HzRGs in the near-IR regime, an overlap between both samples exists. \citet{Norris2011} suggest that IFRS might also follow a correlation between near-IR flux density and redshift. This suggestion has been supported by \citet{Collier2014} and \citet{Herzog2014} who find that those IFRS with spectroscopic redshifts are consistent with this suggested correlation. Similarly, ultra-steep radio spectra~($\alpha \lesssim -1.0$) are known to be successful tracers of high-redshift galaxies~(e.g.\ \citealp{Tielens1979,McCarthy1991,Roettgering1994}). The classes of HzRGs and IFRS were both found to have steep radio spectra~\citep{Middelberg2011}.\par Studying IFRS in the optical and IR regime is challenging because of their faintness at these frequencies. In contrast, IFRS are relatively bright in the radio regime, making detailed radio studies feasible. Since the radio emission of RL galaxies is dominated by the AGN, radio studies of IFRS can provide insights into the characteristics of the active nucleus, e.g.\ its age.\par This paper aims at studying the broad radio SEDs of IFRS, spanning a frequency range of more than two orders of magnitude. In Sect.~\ref{data}, we present our sample of 34~IFRS from the ATLAS fields and describe the available data for the ELAIS-S1 and CDFS fields which includes the first data on IFRS below 610\,MHz and above 8.6\,GHz. Among others, we are using data of two of the new-generation radio telescopes and Square Kilometre Array~(SKA; \citealp{Dewdney2009}) precursors, Murchison Widefield Array~(MWA; \citealp{Lonsdale2009,Tingay2013}) and Australian Square Kilometre Array Pathfinder~(ASKAP; \citealp{Johnston2007,Johnston2008,DeBoer2009}). We also describe the Plateau de Bure Interferometer~(PdBI) observations---the highest-frequency radio observations of an IFRS so far---and ancillary data of one IFRS in the xFLS field. Based on the available data, we build and fit radio SEDs for the IFRS in the ATLAS fields in Sect.~\ref{building_fitting_radioSEDs} and analyse them with respect to spectral index, turnover, and high-frequency behaviour in Sect.~\ref{discussion_radioSEDs}. In Sect.~\ref{IFRS_xFLS478}, we present a multi-wavelength and radio SED modelling for the IFRS observed with the PdBI. Our results are summarised in Sect.~\ref{conclusion}. The photometric data obtained in Sect.~\ref{data} are summarised in Appendix~\ref{datasection}. Throughout this paper, we use flat $\Lambda$CDM cosmological parameters $\Omega_\Lambda = 0.7$, $\Omega_\textrm{M} = 0.3$, $H_0 = 70$~km~s$^{-1}$~Mpc$^{-1}$, and the calculator by \citet{Wright2006}. The linear scale in $\Lambda$CDM cosmology is limited in the redshift range $0.5\leq z \leq 12$ between 4\,kpc/arcsec and 8.5\,kpc/arcsec. Following \citet{Cameron2011_beta}, we calculate $1\sigma$ confidence intervals of binomial population proportions based on the Bayesian approach.\par

\label{conclusion} We built radio SEDs for 34~IFRS in the CDFS and ELAIS-S1 fields, covering the frequency range between 150\,MHz and 34\,GHz. Based on these SEDs, we found the vast majority of IFRS ($74^{+6}_{-9}$\%) to show steep radio SEDs defined by $\alpha < -0.8$. $6^{+7}_{-2}$\% of the IFRS in our sample are classified as USS sources~($\alpha < -1.3$) and are therefore good candidates for high-redshift sources. The sample of IFRS shows statistically significantly steeper radio SEDs than the broader RL~AGN population. The median spectral index in our IFRS sample is $-0.88$. One IFRS was found to show a flattening or upturning radio SED at 20\,GHz, indicating an additional core component and emphasising the importance of considering high-frequency radio data when studying radio sources, irrespective of their redshift. Our finding that IR-fainter IFRS have steeper radio SEDs supports the hypothesis that IR-fainter IFRS are at higher redshifts.\par We found $3^{+6}_{-1}$\% of our sample are GPS sources and $\geq 15^{+8}_{-4}$\% are CSS sources. These numbers are consistent with the general fraction of GPS and CSS sources in the RL AGN population. This finding implies that at least some IFRS are young AGN in the earliest stages of their evolution to powerful and extended FRI/FRII radio galaxies. However, the intrinsic fraction of GPS and CSS sources in the IFRS population might be higher than in the general RL source population if IFRS are at higher redshifts. Generally, IFRS are prototypical for the class of CSS sources because of their steep radio SEDs and their compactness. Our analysis showed that IFRS with an observed peak in their radio SED are IR-fainter than IFRS without a turnover. However, we do not find evidence that the radio flux densities or the radio-to-IR flux density ratios of peaking IFRS differ from those of the non-peaking subsample.\par We also carried out a detailed analysis of the broadband SED of IFRS~xFLS\,478. This source was observed with the PdBI at 105\,GHz and provided the highest-frequency radio detection of an IFRS. The source was found to have a steep radio SED, potentially indicating a turnover at around 150\,MHz. We did not observe an upturn or flattening in the radio SED at high frequencies, indicating that synchrotron emission dominates over thermal dust emission at least down to a rest-frame frequency of 300\,GHz (1\,mm) if the source is at $z\gtrsim 2$.\par Modified SED templates of known galaxies were found to be inconsistent with the multi-wavelength data of xFLS\,478. However, the data are well described by a radio-FIR SED template composed of a star forming galaxy and an RL~AGN at $z=1.1$ which would make this object the IFRS with the lowest known redshift. This model suggests a star formation rate of around 170\,$M_\odot$\,yr$^{-1}$.\par

1607.02707

1607

1607.00228_arXiv.txt

{We study the appearance of the neutron star - accretion disk system as seen by a distant observer in the UV/X-ray domain. The observed intensity spectra are computed assuming non-spherical geometry of the whole system, in which outgoing spectrum is not represented by the flux spectrum, the latter being valid for spherically symmetric objects. Intensity spectra of our model display double bumps in UV/X-ray energy domains. Such structure is caused by the fact that the the source is not spherically symmetric, and the proper integration of intensity over emitted area is needed to reproduce observed spectral shape. Relative normalization of double bump is self consistently computed by our model. X-ray spectra of such a type were often observed in LMXB with accretion disk, ultra luminous X-ray sources, and accreting black hole systems with hot inner compact corona. Our model naturally explains high energy broadening of the disk spectrum observed in some binaries. We attempted to fit our model to X-ray data of XTE~J1709-267 from {\it XMM-Newton}. Unfortunately, the double intensity bump predicted by our model for LMXB is located in soft X-ray domain, uncovered by existing data for this source. } {accretion,accretion disks -- stars:neutron -- X-ray:binaries}

Low mass X-ray binaries (LMXB) are binary systems, where the primary star is a compact object, being the black hole or the neutron star. The secondary is a late type main sequence star, classified as K or M of low mass slightly less than the solar mass, $M_{\odot}$. In many of those sources accretion of matter onto the compact object occurs if the system is tight enough. Accretion can proceed under different scenarios. In principle, when falling matter has non-zero initial angular momentum, then the accretion disk may form down to the innermost marginally stable orbit (ISCO). Since the time when the standard accretion disk model was defined by Shakura \& Sunyaev (1973), theoreticians developed the idea taking into account relativistic corrections (Novikov \& Thorne 1973) and radial advection (Abramowicz et al. 1988). According to all those models, accretion disks in LMXB are hot in their central part, having the gas temperature of the order of $10^{6-7}$ K. For the purpose of this paper, relativistic corrections and radial advection were neglected, since we model basic disk emission. If the compact object in LMXB is a neutron star, then we should expect that the neutron star is hot, with $T_{\rm eff,NS}$ up to a few $10^{6-7.5}$ K. Such system is bright in X-ray domain, and the derivation of the total shape of its intensity spectrum is the principal aim of this paper. In this research project we computed the observed intensity of the whole LMXB system containing a neutron star with the accretion disk. The secondary is not visible in X-rays and that star is only important to estimate the size of the Roche lobe, and therefore, the outer disk radius. Since the emitting source such as LMXB is not spherically symmetric then we have to reject standard formula that the observed luminosity is just proportional to the flux (in ergs per second per one cm$^{2}$) emitted from the unit surface of a spherical star (Mihalas 1978). In case of the disk-like system we have to compute the observed intensity starting from the basic formula i.e. Eq.~[1-27]. in Mihalas (1978). In this paper we derived such formula appropriate to the LMXB with neutron star, but the same calculation can be done for any non-spherical system. We calculated the observed intensity assuming that both the neutron star and the accretion disk radiate as black bodies. The radial effects of mutual attenuation were fully taken into account. Our model spectra seen at different viewing angles display double bumps in the UV/X-ray energy domain. The relative strength of bumps depends on the neutron star and the inner accretion disk effective temperatures. Furthermore, the high energy bump originating from partially attenuated neutron star depends on the viewing angle, which is not the case of the neutron star alone. Such a type of continuum spectra are very often observed in LMXB with accretion disks, ultra luminous X-ray sources (ULXs), and black hole accretion disks with hot corona. Our model should be used for any non-spherical systems containing the accretion disk and the inner emitting source of different temperature, which can be for example, the hot compact corona (Fabian et al. 2015). In Sec.~2 we present the source geometry and derive the observed intensity of the whole system as a function of the viewing angle. The resulting appearance of non-spherical systems with parameters typical for LMXB is drawn in Sec.~3. The first fit of our models to the X-ray data of XTE J1709-267 is shown in Sec.~4. We discuss and conclude our work in Sec.~5.

\label{sec:concl} In this paper we theoretically explained the appearance of non-spherical emission by proper computation of the amount of energy which goes directly towards observer. We defined full set of terms for observed intensity of a non-spherical system, consisting of a neutron star and the accretion disk. We took into account attenuation effects as seen by a distant observer at various aspect angles. We demonstrated that the overall continuum shape shows two peaks. The lower energy peak is caused by the accretion disk emission, whereas higher energy bump is due to the neutron star. The position of the transition between peaks and their relative normalization is self consistently computed by our formulae. Contrary to the spherically symmetric emission the neutron star contribution in our model depends on the aspect angle, due to attenuation effect. In this paper we showed, that if the observed intensity of a source was correctly computed, even without relativistic correction and without light from the boundary layer, then we obtain the double bump source spectrum due to the double non-spherically symmetric, emitting region. Such double bump spectrum, or the broadness of the spectral disk component, are very often observed in X-ray band of accreting objects. We attempted to fit the X-ray spectrum of LMXB XTE~J1709-267, downloaded from {\it XMM-Newton} archive. Nevertheless, the spectral data are too narrow co cover and fit our model, because we have no points below 0.5 keV. The additional points in optical and UV domains are necessary to discriminate our model parameters. Unfortunately, the double bump structure predicted by our model of LMXB with an accretion disk is difficult to observe since it falls into a photon energy range less then 0.5 keV. If the non-spherical emitting system contains hot compact corona (Fabian et al. 2015) instead of the neutron star, the double peak structure should be visible in X-ray domain. We argue that for some parameters the total spectrum from both components can explain soft X-ray excess in those objects. Our model fully predicts the broadness of the disk component observed in some sources (Kolehmainen et al. 2011, and priv. com.). \Acknow{This research was supported by Polish National Science 2013/11/B/ST9/04528, % 2015/17/B/ST9/03422, 2015/18/M/ST9/00541 and by Ministry of Science and Higher Education grant W30/7.PR/2013. It has received funding from the European Union Seventh Framework Program (FP7/2007-2013) under grant agreement No.312789.}

1607.00228

1607

1607.02477_arXiv.txt

We present a general method to identify infalling substructure in discrete datasets with position and line-of-sight velocity data. We exploit the fact that galaxies falling onto a brightest cluster galaxy (BCG) in a virialised cluster, or dwarf satellites falling onto a central galaxy like the Milky Way, follow nearly radial orbits. If the orbits are exactly radial, we show how to find the probability distribution for a satellite's energy, given a tracer density for the satellite population, by solving an Abel integral equation. This is an extension of \citet{Ed16}'s classical formula for the isotropic distribution function. When applied to a system of galaxies, clustering in energy space can then be quantified using the Kullback-Leibler divergence, and groups of objects can be identified which, though separated in the sky, may be falling in on the same orbit. This method is tested using mock data and applied to the satellite galaxy population around M87, the BCG in Virgo, and a number of associations are found which may represent infalling galaxy groups.

\label{sec:introduction} In the hierarchical model of galaxy formation, elliptical galaxies and the stellar haloes of spiral galaxies are built up gradually by prolonged periods of accretion. In this picture, early-type galaxies form during an early phase of dissipational star formation fed by gas-rich mergers, creating compact cores, while their subsequent evolution is mainly driven by abundant minor mergers, feeding the outer regions of the galaxy and bringing them onto the local size-mass relation \citep[e.g.][]{Shen2003, Na09, Hopkins2010}. This phenomenon is even more important in clusters, where the brightest cluster galaxies (BCGs) that reside in their centres are thought to have acquired the majority of their stars through the accumulation and subsequent destruction of satellite galaxies~\citep{La13,Co15}. BCGs have haloes that are embedded in the intracluster light, and it is not obvious whether a useful distinction between the halo and the intra-cluster medium is even possible~\citep{Go05}; the disruption of satellite galaxies as they are funnelled by dynamical friction from the intra-cluster medium onto the central BCG then creates an extensive envelope. Indeed, it is possible that the progenitors of today's BCGs are the extremely compact `red nuggets' that have been observed at redshifts $z \sim 2$ \citep[e.g.][]{Trujillo2007, vanDokkum2008}. The evolution of the BCG is therefore dominated by the addition of stars and globular clusters that have formed outside the BCG itself. In the Local Group, a much lower-mass example compared to galaxy clusters, the biggest members show convincing evidence of a structured satellite galaxy population. For instance, the Large Magellanic Cloud, together with its satellites, may be part of an extended group that is on its first infall onto the Milky Way~\citep{Ko15}, while M31 seemingly has an extended thin disk of satellites~\citep{Ib13} that may be the result of group infall and accretion~\citep[e.g.][]{Bo14}. In these systems, detailed searches for substructure are made possible by the high-quality six-dimensional phase-space data that are available for individual stars \citep[e.g.][]{Xue2011}. Given these findings, the extreme densities found in cluster environments make it very likely that the outer parts of BCGs are also permeated with radially infalling satellite groups. However, substructure identification in these much more distant systems cannot be carried out using the same methods as in the Local Group as it is not possible to resolve individual stars, and bright stellar proxies such as globular clusters, planetary nebulae and satellite galaxies must be used instead. Another difficulty is that these studies must be carried out using projected data, as only three of the full six-dimensional phase space coordinates are usually available -- namely, position on the sky and line-of-sight velocity -- and this gives rise to much larger uncertainty. Nevertheless, a number of recent studies have used globular cluster kinematics to find evidence for recent accretion events \citep[e.g.][]{Cote2003,Schuberth2010,Ro12,Lo15}, while searches for apparent photometric disturbances such as shells and tidal tails have also made strong cases for recent and ongoing accretion \citep{Tal2009}. Extended sheets or pancakes of satellite galaxies have also been tentatively identified in the outskirts of clusters~\citep{Fa13}. Here, we introduce a new method to identify members of the same infalling satellite galaxy group in a cluster, using only projected galactocentric distances and line-of-sight velocities. We argue that, in the outer parts of clusters, galaxies are falling in on almost radial orbits. This suggests an appealing simple ansatz, namely that the distribution function (DF) is \begin{equation} F \propto \delta(v_\theta)\delta(v_\phi) f(E) \label{eq:dfradial} \end{equation} where $v_r, v_\theta,v_\phi$ are velocity components resolved with respect to the spherical polar coordinates and $E$ is the energy. In Section 2, we show that this leads to an Abel inversion, which can be performed exactly. In other words, for any tracer density of objects, a DF of this form can be found, although -- as always -- we must check for positive definiteness a posteriori. As an application, Section 3 applies the method to the dataset of satellite galaxies around M87, one of the most massive galaxies in the local universe. This giant E0 elliptical resides at the centre of the Virgo cluster, and its environment has been catalogued extensively by \citet{Binggeli1985} and \citet{Ki11}. Here, we use the carefully-selected subsample of Virgo galaxies considered to be certain M87 satellites, compiled by \citet{Ol15}, and identify possible substructures by looking for objects which are clustered in energy space, and hence falling onto M87 on the same orbital path but with different phases. Finally, we summarise in Section 4 and consider possible extensions, applications and future prospects for our method.

Strong tests of the current cosmological paradigm are provided by the abundance of substructure. As larger structures are assembled hierarchically from mergers and accretion, we should be able to identify fossil signatures of these events in galaxies and clusters. In particular, galaxy clusters are characterised by a virialised region within which all components -- galaxies and dark matter -- are in dynamical equilibrium surrounded by infall zones in which groups of galaxies are falling into the relaxed cluster. The identification of substructure has mainly been studied in the context of the Milky Way halo. Here, the existence of (in the best cases) six dimensional phase space coordinates makes the problem easier. For example, the use of the actions and frequencies~\citep{Mc08,Sm09} has been advocated to identify past merger events in the Milky Way. The problem of the identification of substructure in projected datasets -- in which only positions on the sky, line-of-sight velocities and heliocentric distances are available -- is harder. It is also of much greater interest and applicability, both to nearby galaxies and to galaxy clusters. Here, we have introduced a new formalism to describe the dynamical state of the outer parts of galaxies and galaxy clusters. The infalling motions of objects are assumed to be generated by purely radial orbits. This means that the probability distributions of observable quantities can be inferred, given the density of the infalling tracers and an estimate of the gravitational potential in which they move. We have provided a general algorithm to do this for spherical potentials. This enables us to search for infalling groups of objects, even though they may be scattered across the sky. Algorithms to quantify substructure in projected data are scarce. The only other one known to us looks for shells, which can be quantified by the characteristic ``chevron pattern'' discernible in line-of-sight velocity and position plots~\citep{Ro12}. As a practical application of our method, we have examined the dataset of satellite galaxies around M87. The extended envelope of M87 has been built from a deluge of smaller satellite galaxies, which may have accreted along preferred directions. Hence, we expect correlations in the satellite galaxy dataset, as some of the satellites may have fallen in along the same orbital path. Our algorithm exploits these correlations to identify kinematically similar substructure. In the case of M87, we have identified a number of possible galaxy associations. These are satellite galaxies whose position and kinematics are consistent with infall on the same radial path. This is expected in theories like $\Lambda$CDM in which the infall of satellites is coherent rather than random. This provides proof of principle that our algorithm can be applied to real data to extract useful results. A possible test of the galaxy associations around M87 may be afforded by deep photometry of the candidates to find the position angles of the major axis. Infalling objects are expected to be stretched out along the orbital path, and so -- if they lie on the orbits conjectured in this paper -- they will be radially distended and their major axes will point towards M87. This effect is seen in the accretion of subhaloes in dissipationless simulations~\citep{Ku07, Ba15} and persists with the inroduction of baryons~\citep{Kn10}. Radial alignment has also been detected observationally in galaxy clusters and groups~\citep[see e.g.,][]{Ha75, Ag06}, though the magnitude of this effect is unclear. However, we would predict significant isophotal alignment of the major axis of our candidates with the cluster radial direction if our associations are real. Another natural application is to globular cluster and satellite galaxy datasets in other nearby galaxies. There are intense observational efforts focussing on completing the surveys of stellar streams and substructure around the Milky Way; however, a complete picture can only be obtained by studying a wider sample of galaxies at greater distances, although this is a much harder problem observationally. A good place to start would be with the satellites and clusters of M31, where coherent streams are readily visible; slightly further afield, the Centaurus group may also be a good candidate. We also note that M87 has a very large number of globular clusters \citep{Ol15a} which may provide additional insights into the accretion history of the galaxy we have studied in this work. We anticipate that this algorithm will be a valuable tool in helping to investigate the build up of structure in the Local Group and beyond. It is surprising that DFs built from only radial orbits have not received much more attention. Perhaps this is because for the fully self-consistent problem (in which the density generates the potential), such DFs fall foul of the radial orbit instability~\citep{Fr84}. However, this objection does not apply to tracer populations, which are moving in an external potential provided largely by other stellar and dark matter populations. We have shown that the radial orbit DF can always be found by Abel transforms via an inversion similar to \citet{Ed16}'s classical work for the isotropic DF. In fact, radial orbit DFs are applicable to a wide range of astrophysical problems. In this paper, we have concentrated on material infalling onto BCGs, but the DFs are also applicable to populations expelled from central nuclei. The hypervelocity stars in the Milky Way are believed to be ejected by the central black hole with speeds from a few hundred to a few thousand kms$^{-1}$. The runaway stars are formed when one component of binary receives a kick as its companion explodes as a supernovae. Both hypervelocity and runaway stars are ejected from the central parts of the Milky Way with such high velocities that they move on almost radial orbits, as shown by simulations of their space motion by ~\citet{Ke14}. Our DFs should have a ready application to the descriptions of these radially ejected populations as well.

1607.02477

1607

1607.03701_arXiv.txt

Ionization-recombination balance in dense interstellar and circumstellar environments is a key factor for a variety of important physical processes, such as chemical reactions, dust charging and coagulation, coupling of the gas with magnetic field and development of instabilities in protoplanetary disks. We determine a critical gas density above which the recombination of electrons and ions on the grain surface dominates over the gas-phase recombination. For this regime, we present a self-consistent analytical model which allows us to exactly calculate abundances of charged species in dusty gas, without making assumptions on the grain charge distribution. To demonstrate the importance of the proposed approach, we check whether the conventional approximation of low grain charges is valid for typical protoplanetary disks, and discuss the implications for dust coagulation and development of the ``dead zone'' in the disk. The presented model is applicable for arbitrary grain-size distributions and, for given dust properties and conditions of the disk, has only one free parameter -- the effective mass of the ions, shown to have a low effect on results. The model can be easily included in numerical simulations following the dust evolution in dense molecular clouds and protoplanetary disks.

\label{intro} An accurate calculation of ionization-recombination balance in dense protoplanetary conditions is essential for understanding various fundamental problems, such as coupling of the gas with magnetic field \citep{Li2014}, accretion processes \citep{Turner2014}, chemistry \citep{Semenov2004,Larsson2012} and dust evolution~\citep{Okuzumi2011b,Akimkin2015}. Both the ionization and recombination processes can arise from several sources. While the treatment of ionization, despite the variety of ionization sources, could be reduced to a single (total) ionization rate, the description of recombination is less straightforward. At sufficiently high densities, the dominant sink of free electrons and ions are dust grains, and the recombination rate non-trivially depends on properties of the grains. Furthermore, collection of electrons and ions leads to non-zero grain charges, which effectively changes the grain-grain \citep{Okuzumi2009} and ion-grain \citep{Weingartner1999} interactions as well as the grain dynamics. Depletion of electrons caused by the presence of dust grains significantly reduces the degree of ionization in dense interstellar conditions \citep{Umebayashi1983,Umebayashi1990,Nishi1991}: In comparison with dust-free gas, the electron-to-ion ratio may drop by as much as a square root of the effective ion-to-electron mass ratio (which is a factor of 74 for a plasma with dominant H$_3^+$ ions, or 231 for N$_2$H$^+$/HCO$^+$ ions). As the ionization controls the coupling of the gas to the magnetic field, and hence the development of the magnetorotational instability \citep[MRI, e.g.,][]{Velikhov1959,Balbus1991}, dust is the essential ingredient for any MRI model. It has been shown that the grain size critically affects the size of a disk's ``dead zone'' \citep{Sano2000,Salmeron2008,Bai2011a,Bai2011b,Dudorov2014}. Nevertheless, analysis of MRI has been usually carried out assuming that properties of dust are fixed. In dense protoplanetary environments, the coagulation of sub-$\mu$m interstellar dust particles becomes an important process. The planet formation in protoplanetary disks requires the dust to form larger and larger aggregates, until gravitational forces become dominant \citep[e.g.,][and references therein]{Testi2014}. There is a clear evidence of grain growth to millimeter and centimeter sizes within protoplanetary disks \citep[e.g.,][]{Perez2015,vanderMarel2015}. However, significant difficulties are found during this coagulation process, such as bouncing barriers \citep[e.g.,][]{Zsom2010} and particle fragmentation \citep[][]{Birnstiel2012} after initial grain compaction and growth. Many theoretical and laboratory studies have greatly advanced our understanding of grain growth and planetesimal formation in recent years \citep[][and references therein]{Dominik2007,Johansen2014}, with particular attention dedicated to dust traps, now detected with ALMA toward protoplanetary disks \citep[][]{vanderMarel2013,vanderMarel2015,Pinilla2015,Flock2015,Zhang2016,Ruge2016}. In dust traps, particles are expected to grow more easily due to the locally enhanced dust-to-gas mass ratio \citep[][]{Booth2016,Surville2016}, although the details of this coagulation process are far from being understood, considering the largely unknown dust properties. As has been already pointed out, the ionization does not only determine dynamical and chemical processes occurring in protoplanetary disks, but also leads to the dust charging and thus affects the coagulation. Collection of electrons and ions results in (on average) negative grain charges due to higher electron velocities. Recently, it has been shown that the coagulation of larger aggregates in protoplanetary disks can be inhibited due to growing Coulomb repulsion between them~-- the resulting electrostatic potential barrier is roughly proportional to the aggregate size \citep[][]{Okuzumi2009,Okuzumi2011a,Okuzumi2011b}. Along with the plasma charging, other charging mechanisms can operate in protoplanetary disks. In \cite{Akimkin2015} the photoelectric emission from grains, induced by stellar radiation and leading to their positive charging, was considered as a mechanism to overcome the electrostatic barrier in upper disk regions. A similar mechanism~-- photoelectric charging due to H$_2$ fluorescence induced by cosmic rays (CRs) -- operates in much deeper regions at the disk periphery \citep{Ivlev2015b}. However, both mechanisms become negligible in dense regions of the disk. We notice that the (still poorly investigated) effect of charging on the dust evolution has recently received increased interest \citep{Carballido2016}. It is noteworthy to mention that the coagulation in protoplanetary disks is accompanied by the formation of porous aggregates characterized by an open, fluffy structure \citep[e.g.,][]{Dominik2007,Okuzumi2009b}. The porosity has been pointed out to have a strong impact on the ionization in protoplanetary disks \citep[][]{Okuzumi2009,Dzyurkevich2013,Mori2016}. However, while the growth of {\it uncharged} aggregates is well studied, an accurate description of their charging as well as of the charging feedback on their further growth poses a serious problem. One of the fundamental difficulties is that, unlike compact spherical grains \citep[whose charging is described using the Orbital Motion Limited (OML) approximation, e.g.,][]{Whipple1981,Fortov2005}, no accurate approximation is known for the electron and ion collection by irregular fluffy aggregates. Given these difficulties, here we leave completely aside in-depth discussion of the porosity effects. In this paper, we present an analytical model which becomes exact in sufficiently dense astrophysical environments and allows us to self-consistently calculate densities of the charged species, in particular -- to obtain the dust charges for arbitrary grain-size distributions. Unlike other known approaches \citep[][]{Ilgner2006,Okuzumi2009,Fujii2011,Dzyurkevich2013,Mori2016}, our model does not make assumptions on the form of the charge distribution, and yields closed analytical expressions for important limiting cases. The latter enables convenient analysis of results in a general form, in terms of a few dimensionless numbers. The presented model has only one free parameter (the effective mass of the ions, which we show to have a low effect on results), and can be easily included in numerical simulations following the dust evolution in dense molecular clouds and protoplanetary disks. We employ the model to verify whether the broadly used approximation of low grain charges is valid for typical protoplanetary disks. Furthermore, we identify a ``dust-dust'' plasma regime, where the grain charge distribution becomes quasi-symmetric with respect to uncharged state. This leads to removal of the repulsive electrostatic barrier and opens a ``coagulation window'' for large aggregates, operating in the inner dense region of protoplanetary disks. Also, we discuss the importance of self-consistent analysis of the ionization and the grain evolution, as there processes are mutually coupled via several mechanisms operating in the disks. The paper is organized as follows. In Section~\ref{dense} we consider the overall ionization-recombination balance and introduce a recombination threshold -- the gas density above which the electron-ion recombination is dominated by the processes on the dust surface. In Section~\ref{charge} we present the grain charge distribution determined by collection of electrons and ions, and point out limiting cases of ``big'' and ``small'' grains. In Section~\ref{densities} we derive the governing equations for the dust-phase recombination regime, complemented with the grain charge distribution, which allow us to calculate densities of the charged species in a general form; to reveal generic properties of the solution, we consider ``monodisperse'' dust (grains of the same size) and investigate separately the big- and small-grain limits. The effect of the grain-size distribution is studied in Section~\ref{size_effect}. We discuss implications of the proposed model for protoplanetary disks in Section~\ref{implications}, and summarize the results in Section~\ref{conclusion}.

\label{conclusion} We have developed an exact analytical model which describes ionization and dust charging in dense protoplanetary disk conditions, for arbitrary grain-size distribution. Unlike previously developed approaches \citep[][]{Ilgner2006,Okuzumi2009,Fujii2011,Dzyurkevich2013,Mori2016}, our model does not make assumptions on the form of the grain charge distribution, and enables convenient analysis of results in a general form, in terms of a few dimensionless numbers, which allows us to identify universality in the behavior of the charged species. The governing equations for different cases are summarized in Appendix~\ref{notations}, Table~\ref{tabsum}. For given dust properties and conditions of the disk, the presented model has only one free parameter (the effective mass of the ions $A$), and is developed for the regime where the dust-phase recombination of free electrons and ions dominates over the gas-phase recombination. A transition to this regime occurs in an electron-ion (EI) plasma (where charged grains still do not play any role in the overall charge neutrality), and is characterized by the dust-phase recombination threshold $(n_{\rm g}/\zeta)_{\rm rec}$ for the gas density. At higher gas densities, $n_{\rm g}/\zeta\gtrsim(n_{\rm g}/\zeta)_{\rm rec}$, charged grains play an increasingly important role in the charge neutrality. We have determined two characteristic parameters, the electron depletion threshold $(n_{\rm g}/\zeta)_{\rm dep}\gg (n_{\rm g}/\zeta)_{\rm rec}$ and the asymptotic threshold $(n_{\rm g}/\zeta)_{\rm asy}\gg(n_{\rm g}/\zeta)_{\rm dep}$, marking, respectively, transitions from the EI to dust-ion (DI) plasma state, and then to the dust-dust (DD) state. The thresholds are determined in such a way that at $n_{\rm g}/\zeta=(n_{\rm g}/\zeta)_{\rm dep}$ electrons and negative grains equally contribute to the total negative charge, while at $n_{\rm g}/\zeta=(n_{\rm g}/\zeta)_{\rm asy}$ ions and positive grains provide equal contribution to the total positive charge. The immediate important implications of the derived results for protoplanetary disks are as follows: \begin{enumerate} \item Unless the dust size distribution is dominated by grains much smaller than $\sim1~\mu$m, larger grains are typically multiply (negatively) charged. In this case, irrespective of the location in the disk, the average grain charge scales linearly with the size. As the size distribution in protoplanetary disk conditions develops towards bigger grains, the conventional approximation of low grain charges may only be used for (very) initial stages of the disk evolution. The presented results are obtained assuming compact dust, and therefore the implication for porous aggregates (which are expected to carry higher charges due to bigger effective sizes) is even stronger. In situations where an aggregate is approximated by a sphere, the effects of porosity can be straightforwardly included by adopting a fractal scaling law $m_{\rm d}\propto a^D$, relating the dust mass and the effective size \citep[with appropriate fractal dimensionality $D<3$, which is known to vary during the dust evolution,][]{Okuzumi2009b}. \item The asymptotic transition to a DD plasma implies that the grain charge distribution becomes quasi-symmetric with respect to $Z=0$. This completely removes the repulsive electrostatic barrier and opens a ``coagulation window'' for large aggregates, operating in the inner dense region of protoplanetary disks. On the other hand, the DD plasma state only ``delays'' the formation of the barrier: the coagulation itself leads to decreasing dust number density and a gradual shift back to DI/EI plasmas. The (re)appearance of the electrostatic barrier in this case, with the maximum achieved in the EI state (where the energy of the barrier normalized to the thermal energy of grains, $\sim |\langle Z\rangle|\Psi_{\rm EI}$, is always very large), may completely inhibit further dust growth. The effect of the barrier can be reduced by various feedback mechanisms operating in the disk and leading to increased local gas-to-dust ratio, such as the dust trapping or moderate fragmentation. \end{enumerate} We point out that the dust evolution, change in the charged species abundances, and development of MRI are strongly interrelated processes whose mutual effect is poorly understood: The dust coagulation increases the ionization degree which, in turn, leads to higher grain charges and prevents further coagulation due to growing electrostatic barrier. On the other hand, higher ionization favors the development of MRI, making a disk turbulent; a moderate turbulence facilitates dust growth by increasing the relative grain velocities, while strong turbulence leads to dust fragmentation. The latter generates small grains, which may decrease the electron fraction (asymptotically, by a factor of $\sqrt{m_{\rm i}/m_{\rm e}}\sim10^2$) and lead to MRI quenching. A complex interplay of these nonlinear processes suggests the existence of multiple positive and negative feedback loops that may dramatically affect the ultimate dust evolution. A rigorous treatment of the ionization fraction and dust evolution could be critical during all stages in the process that links molecular clouds to stellar systems -- this is the motivation behind our work. The exact analytical model presented here can be easily implemented in non-ideal MHD simulations, to properly follow the ionization fraction and the dust growth during the process of protoplanetary disk formation \citep[e.g.,][]{Zhao2016} and evolution \citep[e.g.,][and references therein]{Dullemond2007,Armitage2011}. We provide a FORTRAN source code, applicable for arbitrary dust size distributions and also for porous grains, to calculate abundances of the charged species: http://www.inasan.ru/\~{}akimkin/codes.html

1607.03701

1607

1607.03537_arXiv.txt

We compile an updated list of 38 measurements of the Hubble parameter $H(z)$ between redshifts $0.07 \leq z \leq 2.36$ and use them to place constraints on model parameters of constant and time-varying dark energy cosmological models, both spatially flat and curved. We use five models to measure the redshift of the cosmological deceleration-acceleration transition, $z_{\rm da}$, from these $H(z)$ data. Within the error bars, the measured $z_{\rm da}$ are insensitive to the model used, depending only on the value assumed for the Hubble constant $H_0$. The weighted mean of our measurements is $z_{\rm da} = 0.72 \pm 0.05\ (0.84 \pm 0.03)$ for $H_0 = 68 \pm 2.8\ (73.24 \pm 1.74)$ km s$^{-1}$ Mpc$^{-1}$ and should provide a reasonably model-independent estimate of this cosmological parameter. The $H(z)$ data are consistent with the standard spatially-flat $\Lambda$CDM cosmological model but do not rule out non-flat models or dynamical dark energy models.

In the standard scenario the currently accelerating cosmological expansion is a consequence of dark energy dominating the current cosmological energy budget; at earlier times non-relativistic (cold dark and baryonic) matter dominated the energy budget and powered the decelerating cosmological expansion.\footnote{For reviews of this picture, as well as of the alternate modified gravity scenario, see \cite{RatraVogeley2008}, \cite{Weinbergetal2013}, \cite{Martin2012}, \cite{Joyceetal2016}, and references therein.} Initial quantitative observational support for this picture came from ``lower'' redshift Type Ia supernova (SNIa) apparent magnitude observations and ``higher'' redshift cosmic microwave background (CMB) anisotropy measurements. More recently, cosmic chronometric and baryon acoustic oscillation (BAO) techniques \citep[see, e.g.,][]{Simonetal2005, Morescoetal2012, Buscaetal2013} have resulted in the measurement of the cosmological expansion rate or Hubble parameter, $H(z)$, from the present epoch back to a redshift $z$ exceeding 2, higher than currently probed by SNIa observations. This has resulted in the first mapping out of the cosmological deceleration-acceleration transition, the epoch when dark energy took over from non-relativistic matter, and the first measurement of the redshift of this transition \citep[see, e.g.,][]{FarooqRatra2013b, Farooqetal2013b, Morescoetal2016}.\footnote{See \cite{SutherlandRothnie2015} and \cite{MuthukrishnaParkinson2016} for lower limits on this redshift derived using SNIa and other data. For upper limits on the transition redshift see \cite{Ranietal2015}.} $H(z)$ measurements have also been used to constrain some more conventional cosmological parameters, such as the density of dark energy and the density of non-relativistic matter \citep[see, e.g.,][]{SamushiaRatra2006, ChenRatra2011b, FarooqRatra2013a, Akarsuetal2014, ChimentoRicharte2013, GruberLuongo2014, Bambaetal2014, Ferreiraetal2013, Forte2014, Chenetal2015, Dankiewiczetal2014, Capozzielloetal2014, Mengetal2015, GuoZhang2016, MukherjeeBanerjee2016, Alametal2016}, typically providing constraints comparable to or better than those provided by SNIa data, but not as good as those from BAO or CMB anisotropy measurements. More recently, $H(z)$ data has been used to measure the Hubble constant $H_0$ \citep[][]{Verdeetal2014, Chenetal2016a}, with the resulting $H_0$ value being more consistent with recent lower values determined from a median statistics analysis of Huchra's $H_0$ compilation \citep[][]{ChenRatra2011a}, from CMB anisotropy data \citep[][]{Hinshawetal2013, Sieversetal2013, Adeetal2015}, from BAO measurements \citep[][]{Aubourgetal2015, Rossetal2015, LHuillierShafieloo2016}, and from current cosmological data and the standard model of particle physics with only three light neutrino species \citep[see, e.g.,][]{Calabreseetal2012}. In this paper, we put together an updated list of $H(z)$ measurements, compared to that of \cite{FarooqRatra2013b}, and use this compilation to constrain the redshift of the cosmological deceleration-acceleration transition, $z_{\rm da}$, as well as other cosmological parameters. In the $z_{\rm da}$ analysis here we study more models than used by \cite{FarooqRatra2013b} and \cite{Farooqetal2013b}, now also allowing for non-zero spatial curvature in the XCDM parametrization of dynamical dark energy case and in the dynamical dark energy $\phi$CDM model \citep{Pavlovetal2013}. The cosmological parameter constraints derived here are based on more, as well as more recent, $H(z)$ data than were used by \cite{Farooqetal2015} and we also explore a much larger range of parameter space in the non-flat $\phi$CDM model than they did. \begin{table*} \begin{center} \begin{threeparttable} \caption{Hubble parameter versus redshift data} \begin{tabular}{cccc} \hline\hline \multirow{2}{*}{~~$z$} & ~~$H(z)$ &~~~~~~~ $\sigma_{H}$ &\multirow{2}{*}{~~ Reference\tnote{a}}\\ ~~~~~ & (km s$^{-1}$ Mpc $^{-1}$) &~~~~~~~ (km s$^{-1}$ Mpc $^{-1}$)& \\ \tableline\\*[-4pt] 0.070&~~ 69&~~~~~~~ 19.6&~~ 5\\ 0.090&~~ 69&~~~~~~~ 12&~~ 1\\ 0.120&~~ 68.6&~~~~~~ 26.2&~~ 5\\ 0.170&~~ 83&~~~~~~~ 8&~~ 1\\ 0.179&~~ 75&~~~~~~~ 4&~~ 3\\ 0.199&~~ 75&~~~~~~~ 5&~~ 3\\ 0.200&~~ 72.9&~~~~~~ 29.6&~~ 5\\ 0.270&~~ 77&~~~~~~~ 14&~~ 1\\ 0.280&~~ 88.8&~~~~~~ 36.6&~~ 5\\ 0.352&~~ 83&~~~~~~~ 14&~~ 3\\ 0.380&~~ 81.5&~~~~~ 1.9&~~ 10\\ 0.3802&~~ 83&~~~~~ 13.5&~~ 9\\ 0.400&~~ 95&~~~~~~~ 17&~~ 1\\ 0.4004&~~ 77&~~~~~ 10.2&~~ 9\\ 0.4247&~~ 87.1&~~~~~ 11.2&~~ 9\\ 0.440&~~ 82.6&~~~~~ 7.8&~~ 4\\ 0.4497&~~ 92.8&~~~~~ 12.9&~~ 9\\ 0.4783&~~ 80.9&~~~~~ 9&~~ 9\\ 0.480&~~ 97&~~~~~~~ 62&~~ 2\\ 0.510&~~ 90.4&~~~~~ 1.9&~~ 10\\ 0.593&~~ 104&~~~~~~ 13&~~ 3\\ 0.600&~~ 87.9&~~~~~ 6.1&~~ 4\\ 0.610&~~ 97.3&~~~~~ 2.1&~~ 10\\ 0.680&~~ 92&~~~~~~~ 8&~~ 3\\ 0.730&~~ 97.3&~~~~~ 7&~~ 4\\ 0.781&~~ 105&~~~~~~ 12&~~ 3\\ 0.875&~~ 125&~~~~~~ 17&~~ 3\\ 0.880&~~ 90&~~~~~~~ 40&~~ 2\\ 0.900&~~ 117&~~~~~~ 23&~~ 1\\ 1.037&~~ 154&~~~~~~ 20&~~ 3\\ 1.300&~~ 168&~~~~~~ 17&~~ 1\\ 1.363&~~ 160&~~~~~~ 33.6&~~ 8\\ 1.430&~~ 177&~~~~~~ 18&~~ 1\\ 1.530&~~ 140&~~~~~~~ 14&~~ 1\\ 1.750&~~ 202&~~~~~~~ 40&~~ 1\\ 1.965&~~ 186.5&~~~~~ 50.4&~~ 8\\ 2.340&~~ 222&~~~~~~~ 7&~~ 7\\ 2.360&~~ 226&~~~~~~~ 8&~~ 6\\ [2pt] \hline\hline \label{table:Hzdata} \end{tabular} \begin{tablenotes} \item[a]{Reference numbers: 1.\ \cite{Simonetal2005}, 2.\ \cite{Sternetal2010}, 3.\ \cite{Morescoetal2012}, 4.\ \cite{Blakeetal2012}, 5.\ \cite{Zhang2012} 6.\ \cite{FontRiberaetal2014}, 7.\ \cite{Delubacetal2015}, 8.\ \cite{Moresco2015}, 9.\ \cite{Morescoetal2016}, 10.\ \cite{Alam2016}. } \end{tablenotes} \end{threeparttable} \end{center} \end{table*} We find, from the likelihood analyses, that the $z_{\rm da}$ values measured from the $H(z)$ data agree within the error bars in all five models. They, however, depend more sensitively on the value of $H_0$ assumed in the analysis. These results are consistent with those found in \cite{FarooqRatra2013b} and \cite{Farooqetal2013b}. In addition, the binned $H(z)$ data in redshift space show qualitative visual evidence for the deceleration-acceleration transition, independent of how they are binned provided the bins are narrow enough, in agreement with that originally found by \cite{Farooqetal2013b}. Given that the measured $z_{\rm da}$ are relatively model independent, it is not unreasonable to average the measured values to determine a reasonable summary estimate. We find, for a weighted mean estimate, $z_{\rm da} = 0.72 \pm 0.05\ (0.84 \pm 0.03)$ if we assume $H_0 = 68 \pm 2.8\ (73.24 \pm 1.74)$ km s$^{-1}$ Mpc$^{-1}$. The constraints on the more conventional cosmological parameters, such as the density of dark energy, derived from the likelihood analysis of the $H(z)$ data here, indicate that these data are quite consistent with the spatially-flat $\Lambda$CDM model, the standard model of cosmology where the cosmological constant $\Lambda$ is the dark energy. These $H(z)$ data, however, do not rule out the possibility of dynamical dark energy or space curvature, especially when included simultaneously, in agreement with the conclusions of \cite{Farooqetal2015}. Currently available SNIa, BAO, growth factor, CMB anisotropy, and other data can tighten the constraints on these parameters, and it will be interesting to study these data sets in conjunction with the $H(z)$ data we have compiled here, but this is beyond the scope of our paper. Near-future data will also result in interesting limits \citep[see, e.g.,][]{Podariuetal2001a, Pavlovetal2012, Basseetal2014, Santosetal2013}. The outline of our paper is as follows. In the next section, we discuss and tabulate our new $H(z)$ data compilation. In Sec.\ 3 we summarize how we bin the $H(z)$ data in redshift space and list binned $H(z)$ data. Section 4 summarizes the cosmological models we consider. In Sec.\ 5 we discuss how we compute and measure the deceleration-acceleration transition redshift and tabulate numerical values of $z_{\rm da}$ determined from the $H(z)$ measurements. Section 6 presents the constraints on cosmological parameters, and we conclude in the last section.

From the new list of $H(z)$ data we have compiled, we find evidence for the cosmological deceleration-acceleration transition to have taken place at a redshift $z_{\rm da} = 0.72 \pm 0.05\ (0.84 \pm 0.03)$, depending on the value of $H_0 = 68 \pm 2.8\ (73.24 \pm 1.74)$ km s$^{-1}$ Mpc$^{-1}$, but otherwise only mildly dependent on other cosmological parameters. In addition, the binned $H(z)$ data in redshift space show qualitative visual evidence for the deceleration-acceleration transition, independent of how they are binned provided the bins are narrow enough, in agreement with that originally found by \cite{Farooqetal2013b}. These $H(z)$ data are consistent with the standard spatially-flat $\Lambda$CDM cosmological model but do not rule out non-zero space curvature or dynamical dark energy, especially in models that allow for both. Other data, such as currently available SNIa, BAO, growth factor, or CMB anisotropy data can tighten the constraints on these parameters \cite[see, e.g.][]{Farooqetal2015}, and it is of interest to study how the other data constrains parameters when used in conjunction with the $H(z)$ data we have compiled here.

1607.03537

1607

1607.01829_arXiv.txt

We present a spectroscopic investigation of the Be+sdO binary system HR~2142 that is based upon large sets of ultraviolet observations from the {\it International Ultraviolet Explorer} and ground-based H$\alpha$ observations. We measured radial velocities for the Be star component from these spectra, and computed a revised orbit. In order to search for the spectral signature of the hot subdwarf, we cross-correlated the short wavelength end of each {\it IUE} spectrum with a model hot star spectrum, and then we used the predicted Doppler shifts of the subdwarf to shift-and-add all the cross-correlation functions to the frame of the subdwarf. This merged function shows the weak signal from the spectral lines of the hot star, and a best fit is obtained with a mass ratio $M_2/M_1 = 0.07 \pm 0.02$, companion temperature $T_{\rm eff} \geq 43\pm5$~kK, projected rotational velocity $V\sin i < 30$ km~s$^{-1}$, and a monochromatic flux ratio near 1170 \AA ~of $f_2/f_1 > 0.009 \pm 0.001$. This hot subdwarf creates a one-armed spiral, tidal wake in the disk of the Be star, and we present a circumbinary disk model that can explain the occurrence of shell absorption lines by gas enhancements that occur where gas crossing the gap created by the subdwarf strikes the disk boundaries. The faint companion of HR~2142 may be representative of a significant fraction of Be stars with undetected {\bf former} mass donor companion stars.

% Close binary stars are progressively more common among stars of higher mass \citep{duchene2013}, so we expect that binary interactions will play a fundamental role in the evolution of a significant fraction of massive stars \citep{demink2014}. Roche lobe overflow (RLOF) in close binaries with extreme mass ratios may lead to mergers or a common envelope stage, but in systems with comparable mass stars, RLOF can proceed in a stable manner in which the mass donor star eventually becomes stripped down to a much lower mass \citep{wellstein2001}. RLOF provides mass and angular momentum to the mass gainer star, so that products of binary interaction may appear as a population of rapidly rotating stars \citep{demink2014}. \citet{pols1991} first suggested that such mass gainers might be associated with the rapidly rotating Be stars, emission-line stars surrounded by out-flowing circumstellar disks \citep{rivinius2013}. This idea is supported by the fact that most of the massive X-ray binaries consist of a Be star with an orbiting neutron star companion, the remnant of a massive donor star \citep{reig2011}. However, mass donors below the Chandrasekhar limit will appear as hot subdwarfs or white dwarfs, and such systems with faint companions may account for a large fraction of Be stars \citep{shao2014}. Detection of such faint, subdwarf companions to Be stars is difficult because of the huge flux contrast between the Be star and companion. The stripped down donor stars are expected to be much hotter than the Be stars, so detection is favored in the ultraviolet part of the spectrum where the flux ratio increases. In fact, the first {\bf direct spectroscopic} detection of a hot subdwarf companion of the Be star $\phi$~Per was made by \citet{thaller1995} through investigation of ultraviolet (UV) spectra made with the {\it International Ultraviolet Explorer (IUE)} satellite. The nature of the companion was revealed in subsequent UV spectroscopy from {\it Hubble Space Telescope} \citep{gies1998}, and recent interferometric observations with the CHARA Array have led to an astrometric orbit and masses for the $\phi$~Per system \citep{mourard2015}. Encouraged by the UV detection of the companion of $\phi$~Per, we subsequently launched an investigation of the UV spectra of Be binaries that were frequently observed by {\it IUE}, and our work led to the discovery of the hot subdwarf companions of FY~CMa \citep{peters2008} and 59~Cyg \citep{peters2013}. The presence of a hot companion in $\phi$~Per was first suggested by emission line variations caused by heating of the Be star's disk gas by the hot companion \citep{poeckert1981,hummel2001}, and observations of similar emission line variations may indicate the presence of hot subdwarf companions to the Be stars $o$~Pup \citep{koubsky2012} and HD~161306 \citep{koubsky2014}. Here we turn our attention to the Be binary HR~2142, another system with an excellent set of UV spectra in the {\it IUE} archive. The spectroscopic emission features of HR~2142 (HD~41335, V696~Mon; B1.5 IV-Vnne) were first noted by \citet{campbell1895}, and it has since attracted much attention among spectroscopists because of its strong Balmer emission lines that vary periodically in a predictable manner. \citet{peters1971} observed the development of strong absorption components in the Balmer lines, and their cyclic appearance led \citet{peters1972} to propose that these shell-type absorptions occur during a short part of the orbit in a binary system. Subsequent optical spectroscopy by \citet{peters1983} confirmed the binary nature of HR~2142 through measurement of the radial velocity variations associated with the orbital period of $P=80.86$ days. Early observations with {\it IUE} demonstrated that many transitions in the UV also displayed the shell line variations found in the Balmer lines \citep{peters2001}. \citet{waters1991} examined the spectral energy distribution of HR~2142 in the infrared and showed that the companion cannot be a cool giant (as found in Algol binaries). Instead, they proposed that the companion is a hot subdwarf or ``helium star'', the stripped down remains of the donor star in this interacting binary. Thus, HR~2142 is a prime candidate for detection of a hot companion through a search for its flux contribution in the UV. This paper presents the results of our analysis of the collection of the {\it IUE} spectra of HR~2142. Section 2 describes the UV spectra and a large sample of complementary H$\alpha$ spectroscopy. We present radial velocities of the bright Be star component measured from these spectra, and we calculate an updated radial velocity curve from these data. Section 3 describes our cross-correlation analysis of the far-UV part of the spectrum that provides the first positive evidence of the existence of a hot subdwarf companion. In Section 4, we document the appearance of the shell line variations in the UV and H$\alpha$, and we argue in Section 5 that these shell episodes result from tidal wakes in the disk of the Be star that are caused by the gravitational force of the companion. Our results are summarized in Section 6. Preliminary results of this study appeared earlier in several conference proceedings \citep{peters2001,peters2002,peters2015}.

% Our investigation of the UV and optical spectrum of HR~2142 has led to a new understanding of the nature of this binary system and its circumbinary disk. We used the revised spectroscopic orbit of the Be star to establish the orbital geometry and to predict the radial velocity of the faint companion at the times of observation. These predicted Doppler shifts were applied to cross-correlation functions of the far-UV spectra of HR~2142 with a model spectrum for a hot companion, and the shift-and-added CCFs reveal a faint peak caused by the spectral features of the companion. The companion is a hot and small subdwarf, the stripped down remains of the originally more massive star in this system that transferred both mass and angular momentum to the Be star. Our current results lead to mass estimates of $9 M_\odot$ and $0.7 M_\odot$ for the Be star and hot companion, respectively. Future observations of the orbital motion of the distant tertiary companion of HR~2142 \citep{esa1997,horch2002,oudmaijer2010} will help establish the total system mass. HR~2142 is now the fourth Be+sdO system where UV spectroscopy has led to the detection of the spectral contribution of the hot companion \citep{gies1998,peters2008,peters2013}. The hot subdwarf is small, $R_2/R_\odot > 0.13$, and it has a relatively low luminosity, $\log L/L_\odot > 1.7$. In fact, the companion may have a mass and radius typical of a white dwarf \citep{provencal1998}, although it is much hotter than most. Binary evolutionary models may help determine if the hot companion of HR~2142 is still in a core He burning stage or some more advanced stage (A.\ Schootemeijer et al., in preparation). We have documented the orbital variations of the shell line absorptions in the UV spectrum and H$\alpha$, and we suggest that these are formed mainly in dense gas regions where a tidal wake caused by the companion intersects with boundaries of a gap in the disk of the Be star. Shocks are formed in these regions as gas crossing the gap strikes the dense gas of the disk, and the resulting shock heating helps to explain the presence of the shell components in high ionization species like \ion{Si}{4} and \ion{C}{4}. It may be possible in the future to detect the gap in the disk and the tidal wake regions through high angular resolution observations with optical long baseline interferometry. {\bf Furthermore, the heated tidal wake regions may be responsible for the emission line formation in the vicinity of the companion that is observed in some Be+sdO binaries (in addition to direct heating by the flux of the companion; \citealt{hummel2001}). The \ion{He}{1} $\lambda 6678$ emission line shows evidence of localized heating in the cases of $\phi$ Per \citep{stefl2000}, FY~CMa \citep{peters2008}, and 59~Cyg \citep{peters2013}, so we might expect similar emission variations in the spectrum of HR~2142. There is some evidence of a variable emission component that fills in the wide absorption profile of \ion{He}{1} $\lambda 6678$ in our spectra, but it is much weaker than observed in the other Be+sdO systems, perhaps due to the lower luminosity of the subdwarf. } {\bf It was the remarkable orbital variations in the shell line spectrum that first drew our attention to HR~2142 \citep{peters1971,peters1972}. If the circumbinary disk model of gap crossing gas is applicable to other Be+sdO binaries, then we might naively expect that they should also display such shell line variations in their spectra. However, to our knowledge the appearance of shell lines near Be star superior conjunction has only been reported for the sudden appearance of \ion{Fe}{4} lines in the FUV spectrum of $\phi$~Per \citep{gies1998}. We speculate that tidal wakes may occur in other Be binaries, but that they only present observable shell lines if (1) the Be star's disk is particularly massive and has substantial density out to the vicinity of the companion, (2) the orbital inclination is close to $i=90^\circ$ so that the disk perturbations caused by the tidal wakes reach to a vertical displacement that occults our line of sight to the Be star, and (3) the companion is not so massive that it creates a gap that is too wide to facilitate gas migration across the gap. HR~2142 may represent a fortuitous conjunction in which all three conditions appear to be met. } The example of HR~2142 illustrates the importance of binary Be stars as laboratories for investigations of disk dynamical processes.

1607.01829

1607

1607.05726_arXiv.txt

We present new, more precise measurements of the mass and distance of our Galaxy's central supermassive black hole, Sgr A*. These results stem from a new analysis that more than doubles the time baseline for astrometry of faint stars orbiting Sgr A*, combining two decades of speckle imaging and adaptive optics data. Specifically, we improve our analysis of the speckle images by using information about a star's orbit from the deep adaptive optics data (2005 - 2013) to inform the search for the star in the speckle years (1995 - 2005). When this new analysis technique is combined with the first complete re-reduction of Keck Galactic Center speckle images using speckle holography, we are able to track the short-period star S0-38 (K-band magnitude = 17, orbital period = 19 years) through the speckle years. We use the kinematic measurements from speckle holography and adaptive optics to estimate the orbits of S0-38 and S0-2 and thereby improve our constraints of the mass ($M_{bh}$) and distance ($R_o$) of Sgr A*: $M_{bh} = 4.02\pm0.16\pm0.04\times10^6~M_{\odot}$ and $7.86\pm0.14\pm0.04$ kpc. The uncertainties in $M_{bh}$ and $R_o$ as determined by the combined orbital fit of S0-2 and S0-38 are improved by a factor of 2 and 2.5, respectively, compared to an orbital fit of S0-2 alone and a factor of $\sim$2.5 compared to previous results from stellar orbits. This analysis also limits the extended dark mass within 0.01 pc to less than $0.13\times10^{6}~M_{\odot}$ at 99.7\% confidence, a factor of 3 lower compared to prior work.

\label{intro} Following the motions of stars in the center of our Galaxy has given many insights into the properties of the gravitational potential in which they move. The measurement of the high proper motions, and later, accelerations of these stars implies that they move in the gravitational potential of a concentrated dark mass (\citealt{1997MNRAS.284..576E}; \citealt{1998Ghez}; \citealt{2000Ghez}; \citealt{2002MNRAS.331..917E}). With further observations, these stellar motions have provided strong evidence for the presence of a supermassive black hole (SMBH) at the Galactic Center (Sgr A*) with a mass of about $4 \times 10^{6}$ $M_{\odot}$. Once the star S0-2 went through closest approach in 2002, it was possible to fit its motion with a Keplerian orbit (\citealt{2002Natur.419..694S}; \citealt{2003ApJ...586L.127G}). In addition to the mass of the SMBH, stellar orbits with measured radial velocities have been used to determine the distance to the Galactic Center ($R_o$; \citealt{2003ApJ...586L.127G}; \citealt{2003ApJ...597L.121E}). With S0-2's short orbital period, this star provides the best constraint on the mass of the central black hole and $R_o$ from stellar orbits to date (e.g., \citealt{2008Ghez}, \citealt{2009Gillessen}). The focus of recent work has been to continue assessing the central black hole's properties as well as exploring the potential for using the measurement of stellar motions to test general relativity. The mass of Sgr A* ($M_{bh}$) and the distance to the Galactic Center ($R_{o}$) are both important ways of characterizing this unique region of our Galaxy and putting it in context with others galaxies. Measuring the mass of the central supermassive black hole allows it to be compared to supermassive black holes in the centers of other galaxies. With the mass of Sgr A*, the Milky Way can be added to observed correlations between the mass of the central SMBH and other galactic properties, such as velocity dispersion of stars and bulge luminosity (\citealt{2000ApJ...539L...9F}; \citealt{2002ApJ...574..740T}; see \citealt{2013ARA&A..51..511K} for a review). $R_o$ is a key parameter characterizing our galaxy's size, mass and kinematics. The adopted value of $R_o$ affects estimates of the Milky Way's rotation curve and thereby also measurements of the mass and shape of the dark matter distribution (e.g., \citealt{2000MNRAS.311..361O}). An independent, very accurate measurement of $R_o$ could possibly be used to calibrate stellar distance indicators, such as RR Lyrae and Cepheids, which are important steps on the cosmic distance ladder (see \citealt{1993ARA&A..31..345R}). $R_o$ additionally serves to calibrate the extinction towards the Galactic Center (e.g., \citealt{2010A&A...511A..18S}). The mass-to-distance ratio of Sgr A* as derived from stellar orbits is also necessary to determine the predicted size of the black hole shadow that will be observed by the upcoming Event Horizon Telescope, which can be used as a null hypothesis test of general relativity (see \citealt{2015Psaltis}). Finally, future tests of general relativistic effects on the motion of S0-2 depend on accurate measurements of the gravitational potential due to the SMBH in the Newtonian regime. Until now, our group has used only S0-2 to constrain $R_o$ and the mass of Sgr A*. This is because S0-2 is unique: it is bright (14.2 in the K band) and has a short orbital period (16.2 years). We therefore have been able to track its motion since Keck observations of the Galactic Center began in 1995 such that our observations now cover more than one full orbit of this star. We would like to also use other short-period stars in the Galactic Center to help determine the gravitational potential, but ideally only those stars with high orbital phase coverage like S0-2. It has been shown for visual binary stars that if less than 40 - 50\% of a body's orbit is covered by astrometric observations, the orbital parameters estimated from the data are systematically biased from their true values (\citealt{2014Lucy}). We therefore only use additional short-period stars to constrain the black hole mass and $R_o$ that conservatively have at least this minimum orbital phase coverage. Achieving this minimum orbital coverage for other short-period stars is a challenge because those stars are fainter than S0-2 by more than an order of magnitude, making them difficult to track in the first 10 years of speckle imaging data. In this work, we present a method to increase the orbital phase coverage of these faint stars at the Galactic Center through a complete reanalysis of this data set. In all past analyses of stellar orbital parameters using the speckle data, the data from each observation run has been treated independently. Stars are blindly searched for in the summed image from a given epoch of observation, as if the Galactic Center had never been observed before. No information from other observations is used in this search. In this work, we present a new methodology for analyzing the speckle images that does use information from the much deeper adaptive optics images and the vast improvement that has been made in the knowledge of the central black hole's properties. We use constraints on a star's orbit from the deep adaptive optics data to inform the search for the star in the earlier speckle years. As a pilot study for this new methodology, we apply this technique to S0-38, one of the three stars at the Galactic Center with an orbital period of less than 20 years (in addition to S0-2 and S0-102; see \citealt{Meyer12}), with the ultimate goal of using this star as an additional constraint on the black hole mass and $R_o$. At a magnitude of K=17.0, S0-38 is consistently detected in our deep adaptive optics data taken from 2005 - 2013 but it has not previously been detected in our speckle imaging data taken from 1995 - 2005. S0-38 is an ideal star for the application of this methodology because it is consistently detected in all 21 adaptive optics images and its radial velocity has been measured (\citealt{2009Gillessen} and this work), so its orbit is well known even with just over 40\% of its orbit covered by AO observations. Our results are also made possible by a new reduction of the speckle data using the more sophisticated reconstruction algorithm called speckle holography (\citealt{2013SchoedelHolo}). This is the first work that includes the speckle holography re-reduction of all Keck Galactic Center speckle data. The paper is organized as follows: Section \ref{datasets} describes the data sets used in this work, including the results of the new speckle holography reduction on the full set of Keck Galactic Center speckle imaging data. Section \ref{forcemethod} describes the new methodology of analyzing the speckle images as applied to S0-38. Section \ref{results} contains the results of the S0-38 orbital analysis, including improved constraints on mass and $R_{o}$ as well as extended mass.

\label{discussion} \subsection{Improvements with the Additional Information from S0-38} \label{discussion1} We have demonstrated the power of our new methodology of searching for speckle holography detections of S0-38 using the information in the AO detections by applying it to S0-38. The combination of S0-2 and S0-38 leads to a significant improvement in our knowledge of $R_o$ and the mass of Sgr A* compared to using the information from S0-2 alone. The reasons why the addition of S0-38 improves the constraints on $M_{bh}$ and $R_o$ so much are described here. The results of fitting S0-38 alone show that S0-38 primarily gives information about the position and velocity of Sgr A* on the plane of the sky ($x_o$, $y_o$, $V_x$, and $V_y$). Figure~\ref{fig:MassRo_PosVel} shows the joint probability distributions of these four parameters with $M_{bh}$ and $R_o$ from the fit of S0-2 alone and in the simultaneous fit of S0-2 and S0-38. In the fit of S0-2 alone, $M_{bh}$ and $R_o$ are both correlated with $y_o$, demonstrating that a better constraint on $y_o$ would lead to a better constraint on $M_{bh}$ and $R_o$. In the simultaneous fit, the errors on $y_o$ and thereby $M_{bh}$ and $R_o$ are significantly reduced. So it is specifically through improved knowledge of $y_o$ that S0-38 decreases the errors on $M_{bh}$ and $R_o$ when simultaneously fit with S0-2. The fact that $M_{bh}$ and $R_o$ are not correlated with $V_y$ in Figure~\ref{fig:MassRo_PosVel} additionally shows that the difference in values of this parameter from the separate fits of S0-2 and S0-38 does not affect the $M_{bh}$ and $R_o$ results from the simultaneous fit. \begin{figure}[h] \begin{center} \includegraphics[width=0.5\columnwidth]{massPosVel_S02only_combofit_withjack_nocontour-eps-converted-to.pdf} \includegraphics[width=0.5\columnwidth]{RoPosVel_S02only_combofit_withjack_nocontour-eps-converted-to.pdf} \caption{\label{fig:MassRo_PosVel} Joint probability distributions of $M_{bh}$ (\emph{top}) and $R_o$ (\emph{bottom}) and $x_o$, $y_o$, $V_x$, and $V_y$ as determined by the orbital fit of S0-2 alone (blue dotted line) and by the simultaneous orbital fit of S0-2 and S0-38, with all new speckle holography detections included (red solid line). These probability distributions are shifted by the bias determined in the jack knife analysis (see Appendix \ref{appendix:jackknife}). The correlation between these sets of parameters in the case of S0-2 alone shows that our knowledge of $M_{bh}$ and $R_o$ is limited by how well S0-2 constrains $y_o$. With the addition of the measurements of S0-38's position and RV, $y_o$, and therefore also $M_{bh}$ and $R_o$, are much better determined. The errors in $M_{bh}$ and $R_o$ in the combined S0-2 and S0-38 fit are $\sim$2 and $\sim$2.5 times smaller than in the S0-2-only fit respectively.} \end{center} \end{figure} The reason that the addition of S0-38 gives such an improved constraint on $y_o$ is the orientation of this star's orbit on the plane of the sky. Due to S0-2's orientation on the plane of the sky and the fact that we must omit the detections of S0-2 around its periapse position because of confusion with the NIR counterpart of Sgr A*, S0-2 gives a better constraint on $x_o$ than $y_o$. Figure \ref{fig:XY_S0-2_S0-38} shows the current set of astrometric measurements for S0-2, including those points left out of the fit due to confusion with Sgr A* and other known sources (indicated by open circles), leaving very few unbiased points in the lower third of the ellipse of S0-2's orbit. S0-38 is fortunately on an orbit that is nearly perpendicular on the plane of the sky to that of S0-2, also shown in Figure \ref{fig:XY_S0-2_S0-38}. Therefore, this star can further constrain the $y_o$ position of Sgr A* and in turn the values of $M_{bh}$ and $R_o$. The best-fit solutions of $M_{bh}$ and $R_o$ for each case agree well within 1 sigma. With the addition of the information provided by S0-38, the errors on these parameters decrease by a factor of $\sim$2 and 2.5 respectively. The joint probability distribution of $M_{bh}$ and $R_o$ (Figure \ref{fig:MassRo}) also shows that S0-38 contains different information than S0-2 about these parameters. $M_{bh}$ and $R_o$ are not individually constrained by the orbital fit of S0-38 alone, but their correlation is constrained. The correlation as determined by S0-38 alone has a different slope than the correlation determined by S0-2 alone. This different slope leads to the additional constraints on $M_{bh}$ and $R_o$ when information from S0-38 is added. In the future, the addition of more stars with high orbital phase coverage and orbits perpendicular to S0-2 will lead to even smaller errors on $M_{bh}$ and $R_o$. \subsection{Scientific Implications of New Constraints on $M_{bh}$ and $R_o$} The values of $M_{bh}$ and $R_o$ presented here agree within uncertainties with previous measurements made using stellar orbits. The most recent previous measurements from Keck and VLT data respectively are $4.1\pm0.4$ (\citealt{Meyer12}) and $4.30\pm0.36$ (\citealt{2010RvMP...82.3121G}) million solar masses for the mass of the black hole and $7.7 \pm 0.4$ (\citealt{Meyer12}) and $8.2 \pm 0.34$ (\citealt{2013Gillessen}) kpc for $R_o$. The uncertainties derived from the orbital fit of S0-2 alone are slightly smaller than the uncertainties of these previous measurements due to the increased orbital phase coverage of this star, but the primary improvement in this work is the added information from the orbit of S0-38. The method of using multiple stars' orbits to determine $M_{bh}$ and $R_o$ was also used in \cite{2009Gillessen}. In that work, the orbits of 5 stars in addition to S0-2 were simultaneously fit to determine the gravitational potential parameters. These 5 stars had orbital phase coverages ranging from 12 - 35\% and three of the stars had multiple radial velocity measurements. S0-38 was not included in this set of stars, since it had only been observed from 2005 through 2008 at the time. With our observations now covering over 80\% of S0-38's orbit, just including this one star in addition to S0-2 has significantly decreased the errors on $M_{bh}$ and $R_o$. From the arguments presented in the previous section, this seems to mainly be due to S0-38's orientation on the sky, as well as the more than 80\% orbital phase coverage. Of the 5 other stars used in \cite{2009Gillessen}, two of them have orientations $\sim45$ degrees away from S0-2's vertical orientation, but none has the perpendicular orientation of S0-38. The $R_o$ value presented in this work also agrees with other, recent direct measurements of the distance to the center of our galaxy within $\sim2$~sigma. VLBI measurements of the trigonometric parallax of H$_2$O masers in the star forming region Sgr B2 give a value of $7.9\pm0.8$ kpc (\citealt{2009ApJ...705.1548R}), while dynamical modelling of the nuclear star cluster gives a statistical parallax distance to the Galactic Center of $8.27\pm0.13$ kpc (\citealt{2015MNRAS.447..948C}). Another recent indirect measurement, using VLBI parallaxes and proper motions of over 100 masers to model the motion and structure of the Milky Way, gives a comparable value and uncertainty in $R_o$ to that of the statistical parallax method: $8.34\pm0.16$ kpc (\citealt{2014ApJ...783..130R}). The direct statistical parallax method and indirect modelling method have similar uncertainties in $R_o$, but a value of $\sim2$~sigma higher than presented in this work. It remains to be seen whether continued orbital monitoring and further improvement of $R_o$ constraints from the orbits of S0-2 and S0-38 will maintain this mild disagreement. The direct measurement of $R_o$ presented here has implications for constraints on the structure and kinematics of our galaxy. \cite{2004Reid_propermotion} measured the apparent proper motion of Sgr A*, which is due only to the galactic orbit of the sun if the supermassive black hole is assumed to be at the center of the galaxy. This measurement gave a ratio of the circular rotation speed at the radius of the Sun ($\Theta_o$) and $R_o$ of $29.45 \pm 0.15$ km s$^{-1}$ kpc$^{-1}$. Combining this ratio with the $R_o$ measured in this work gives $\Theta_o = 231$ $\pm 4.3$ km s$^{-1}$. This value of $\Theta_o$ agrees well with the independently measured value found by \citealt{2014ApJ...783..130R} of $240\pm8$ km s$^{-1}$. The new mass value presented here does not significantly change the position of the Milky Way in the observed correlations between mass of the central black hole and host galaxy properties, such as stellar velocity dispersion and mass of the bulge (see \citealt{2013ARA&A..51..511K} for a review). Sgr A*, along with other central black holes in galaxies with pseudobulges, has a lower mass than expected from the tight correlation seen in ellipticals and galaxies with bulges. The new upper limit on the extended dark mass within the orbits of S0-2 and S0-38 is also improved compared to previous work. To compare the limits presented in Section \ref{results} with previous measurements from \cite{2008Ghez} and \cite{2009Gillessen}, we tranform our upper limits to find the extended mass within the apoapse distance of S0-2: 3.1 $\times 10^{11}$ km $=$ 0.01 pc. The 1-sigma (3-sigma) upper limits within this radius are 0.05 (0.14) $\times 10^6 M_{\odot}$ for the S0-2 only fit and 0.04 (0.13) $\times 10^6 M_{\odot}$ for the simultaneous S0-2 and S0-38 fit. \cite{2008Ghez} and \cite{2009Gillessen} found 1-sigma upper limits of 0.12 and 0.17 $\times 10^6 M_{\odot}$ respectively; therefore our new upper limits are a factor of 3-4 lower than previous measurements. S0-38 does improve the limit compared to S0-2 alone, but the main reason for the lower limit compared to earlier work is due to increased time coverage of the orbit of S0-2. Our new limit is still over an order of magnitude more than the $\sim$500 - 1000 $M_{\odot}$ of stellar remnants predicted to be within 0.01 pc (e.g., \citealt{2006ApJ...649...91F}), and other models predict even less mass within 0.01 pc (\citealt{2010ApJ...718..739M}). As observations of the orbits of S0-2 and S0-38 continue, the limits on the extended dark mass within 0.01 pc will continue to decrease. The improved constraints on the gravitational potential that come with the addition of S0-38 also impacts future tests of general relativity in the Galactic Center. General relativistic deviations from pure Keplerian motion are expected to be detectable when S0-2 goes though its time of closest approach to the black hole in 2018. The deviations will be observable as the shift of the measured velocity of S0-2 due to the gravitational redshift. Measuring the deviations from S0-2's Keplerian orbit require as much knowledge of S0-2's Keplerian orbit and the gravitational potential as possible, so the additional constraints from S0-38 are important to this future probe of general relativity. Another observable deviation from a Keplerian orbit is the precession of the point of periapse. This general relativistic precession is confounded by the Newtonian precession due to extended dark mass within stellar orbits. Measuring the general relativistic precession therefore requires the measurement of precession in at least two stars. Additionally, the measurement of the ratio of $M_{bh}$ and $R_o$ from stellar orbits is required to compare the size of the black hole shadow as measured by the upcoming Event Horizon Telescope to theoretical predictions (\citealt{2015Psaltis}). The theoretical half-opening angle of the shadow of Sgr A* as observed from Earth is predicted to be $(5\pm0.2)GM_{bh}/R_{o}c^2$, regardless of the spin of the black hole. With the addition of S0-38 and an increased time baseline of observations, the gravitational radius is now known within $\sim3\%$, which is less than $\pm4\%$ range in the theoretically predicted sizes of the black hole shadow assuming no knowledge of the spin of Sgr A*. In the future, our knowledge of the gravitational potential in the Galactic Center will increase with more observations of S0-2 and S0-38 as well as with the addition of other short-period stars, thereby also increasing the possibility of measuring the effects of general relativity in this extreme environment.

1607.05726

1607

1607.02094.txt

Generation of magneto-optic effects by the interaction of the CMB with cosmic magnetic fields is studied. Effects which generate polarization such as the Cotton-Mouton effect, vacuum polarization and photon-pseudoscalar mixing in external magnetic field are studied. Considering the CMB linearly polarized at decoupling time, it is shown that photon-pseudoscalar mixing in external magnetic field, the Cotton-Mouton effect in plasma and the vacuum polarization in cosmic magnetic field, would generate elliptic polarization of the CMB depending on the photon frequency and magnetic field strength. Among standard magneto-optic effects, the Cotton-Mouton effect in plasma turns out to be the dominant effect in the generation of CMB elliptic polarization in the low frequency part $\nu_0 \sim 10^8-10^9$ Hz with degree of circular polarization $P_C(T_0)\simeq 10^{-10}-10^{-6}$ for magnetic field amplitude $B_{e0}\sim 100\, \text{nG}-1$ nG. The vacuum polarization in magnetic field is the dominant process in the high frequency part $\nu_0\geq 10^{10}$ Hz where the degree of circular polarization at present is $P_C(T_0)\lesssim 10^{-11}$ in the best scenario. The effect of pseudoscalar particles on the CMB polarization is also studied. It is shown that photon-pseudoscalar particle mixing in cosmic magnetic field generates elliptic polarization of the CMB as well and even in the case of initially unpolarized CMB. New limits/constraints on the pseudoscalar parameter space are found. By using current limit on the degree of circular polarization of the CMB, the upper limit of $|g_{\phi\gamma}|<4.29\times 10^{-19}(\text{G}/B_{e0})$ GeV$^{-1}$ for $m_\phi<1.6\times 10^{-14}$ eV in the weak mixing case is found. If $|g_{\phi\gamma}| < 1.17\times 10^{-24}(\text{G}/B_{e0})$ GeV$^{-1}$, a value of the order $|g_{\phi\gamma}|\simeq 10^{-26}(\text{G}/B_{e0})$ GeV$^{-1}$ for $m_\phi\simeq 1.6\times 10^{-14}$ eV in the resonant case, from large scale temperature anisotropy is obtained. Prior decoupling CMB polarization due to pseudoscalar particles is also discussed.

The interaction of light with matter and fields has been intensively studied in the literature and first quantitative studies dates back to Galileo, Newton, Faraday and Maxwell. Among the interesting effects that such interaction represents, there is one class of phenomena which includes the interaction of light (electromagnetic wave) with external electromagnetic fields. These phenomena manifest when an electromagnetic wave propagates through an external electromagnetic field that has been altered by the presence of the incident electromagnetic wave. If there is present only an external electric field, the effects that manifest are called electro-optic effects. Instead, if there is present only an external magnetic field, the effects that manifest belong to the category of the magneto-optic effects. In this work I study only the last effects. Magneto-optic effects not only are important to the established physics but also allow to investigate new effects that have not been found yet. They are generally divided in three main categories that are related to transmission, reflection and absorption of the incident light by the magnetized medium. Depending on the initial polarization of the electromagnetic wave, there are essentially four magnetic-optic effects which belong to the transmission (and not only) category, the Cotton-Mouton (CM) effect, the Faraday effect and two more exotic effects which are the vacuum polarization and the mixing of photons with pseudoscalar (and also scalar) particles in external magnetic field. The reflection category includes essentially only the Kerr effect while the absorption category includes the so called molecular circular dichroism in gases and as will be shown in this work also the photon-pseudoscalar mixing in magnetic field. In the transmission category, the CM effect has been extensively studied in the literature. It has been experimented mostly in gases, liquids, solids and to some degree even in plasma. The CM effect manifest when light propagates in a magnetized medium where the external magnetic field has a transversal component with respect to the direction of light propagation \cite{Born65}. This effect also shares a property with more exotic phenomena such as vacuum polarization (or simply QED effect) and photon-pseudoscalar mixing in magnetic field. All three effects manifest only where there is a transversal component of the external magnetic field with respect to the direction of light propagation. The vacuum polarization has been first proposed and studied in Ref. \cite{Heisenberg:1935qt} and since then has received much attention from both theory and experimental physics. This effect would manifest as a phase shift between the two photon states perpendicular and parallel to the external (transverse) magnetic field that eventually give rise to a birefringence effect which intensity depends on the incident electromagnetic wave frequency. One of the most important achievement from the experimental side, is to measure the acquired QED ellipticity angle of the incident light propagating through the magnetic field. Indeed, this has been the quest for the PVLAS \cite{Bakalov:1994} experiment, BFRT \cite{Cameron:1993mr} experiment and for new generation of experiments \cite{Heinzl:2006xc}. After a first claim of detection of vacuum birefringence by PVLAS experiment \cite{Zavattini:2005tm}, there is still a long way to achieve the required apparatus sensitivity in order to measure the QED predicted ellipticity which is by more than an order of magnitude smaller than the current apparatus sensitivity. At current status, apparatus sensitivity is contaminated with not well understood background noise, must probably from the same apparatus and new methods have also been proposed \cite{Zavattini:2012zs}. The birefringence effect predicted by QED can also be mimicked by another magneto-optical effect, namely the photon-pseudoscalar/scalar mixing in magnetic field. In fact, as it will be shown in this work, mixing of photons with pseudoscalar particles gives rise to both birefringence and dichroism effects. Therefore experiment such as PVLAS, BFRT etc., can in principle find pseudoscalar particles such as axions, ALPs, scalar bosons if the induced birefringence or dichroism signal is bigger than the QED expected signal. Other important experiments that aim to find exotic pseudoscalar particles include the CAST and IAXO experiments \cite{Dafni:2006pc}, ADMX experiment \cite{Asztalos:2009yp} and ALPS-II \cite{Bahre:2013ywa}. Among all magneto-optic effects, the Faraday effect has received much attention in astronomy and cosmology. It manifest when an initial linearly polarized electromagnetic wave interacts with an external magnetic field that has a longitudinal component along the wave propagation direction. This coupling makes possible the rotation of the polarization plane of the incident electromagnetic wave and the rotation angle is proportional to $B_e d$ where $B_e$ is the strength of the external magnetic field and $d$ is the length of the path. Consequently, the Faraday effect has been widely used in radio astronomy as a probe of cosmic magnetic fields, in galaxy clusters and also in the intergalactic space \cite{Kim:1991zzc}. Measurements of the rotation angle of light received from galaxy clusters confirm the presence of a magnetic field inside them, with a magnitude of about few $\mu$G. In the intergalactic space, present studies would suggest a weaker large scale magnetic field with upper limit magnitude $B_e\lesssim 3-1380$ nG, see for example current limits by Planck collaboration \cite{Ade:2015cva} where limits of the order of 1380 nG are set from Faraday effect. On the other hand non observation of gamma rays emission from intergalactic medium due to injection of high energy particles by blazars would suggest a lower value on the strength of extragalactic magnetic field $B_e\geq 10^{-16}-10^{-15}$ G \cite{Neronov:1900zz}. The origin of this field is still unknown and present studies suggest that it may have been created during structure formation or it may have a primordial origin, see Ref. \cite{Grasso:2000wj} for a review on cosmic magnetic fields structure and current updated limits/constraints by Planck collaboration. In this work it is assumed that the magnetic field has a primordial origin and its amplitude is a slowly varying function of space-time coordinates, namely a slowly varying inhomogeneous field in space and time which can be also stochastic in nature, see sec. \ref{sec:7} for details. In connection with the CMB physics, the Faraday effect has been used to probe the existence of primordial magnetic field \cite{Kosowsky:1996yc} present at the decoupling time since it would rotate the polarization plane of the CMB. In fact, it is well known by now that the CMB posses a very small linear polarization that is believed to have been generated at the decoupling time due to Thomson scattering of CMB photons on electrons. Such a polarization is generated because of temperature anisotropies present at the decoupling epoch that eventually generate a position dependent photon intensity on the surface of the last scattering \cite{Bond:1987ub}. Consequently, Thomson scattering of an anisotropic background of photons on electrons would generate linear polarization of the CMB with non zero Stokes parameters $Q$ and $U$ \cite{Kosowsky:1994cy}, \cite{Hu:1997hv}. In general, the linear polarization pattern of the CMB can be decomposed in two modes with opposite parity, the so called E-modes (or gradient modes G) which are the dominant component of the linear polarization and B-modes (or curl modes C) which are the subdominant component of linear polarization, see Refs. \cite{Kamionkowski:1996ks}. The former are generated only by scalar density fluctuations of the cosmological plasma while the latter are generated by either vector or tensor modes. The generation of B-modes is induced by tensor perturbations (gravitational waves) \cite{Crittenden:1993wm}, Faraday rotation of the CMB \cite{Kosowsky:1996yc}, gravitational lensing of the E-mode component \cite{Zaldarriaga:1998ar} and due to primordial magnetic fields \cite{Mack:2001gc} via perturbations sourced by the magnetic field. In general, the spectrum of B-modes is described in multipole moments $l$ of spherical harmonics used to describe linear polarization. The location of the peak signal of B-modes as function of $l$ would give the possibility to distinguish between signals generated by different sources of B-modes. So far, much of attention on the CMB polarization has been focused mostly on the linear polarization. This fact, partially has been influenced by the first experimental observation of E-modes (due to primordial adiabatic scalar fluctuations) by DASI, WMAP and BOOMERANG collaborations \cite{Kovac:2002fg} and also by the fact that many inflationary models predict an almost scale invariant spectrum of gravitational waves, which as already mentioned above, can produce B-modes which are believed to be the `holy grail' of the inflationary theory. Moreover, since Thomson scattering is the most frequent type of scattering in the early universe and because it generates only linear polarization, other types of CMB polarization have been to some extent obscured and the $V$ Stokes parameter has become essentially the `lost along the way' parameter. However, it is well known that light can have two additional types of polarization, circular and elliptic which translate into a nonzero Stokes parameter $V$. After this premise on the CMB linear polarization, several questions come spontaneously. Does the CMB posses only linear polarization? Does it have any degree of circular polarization? If yes, what are the generating mechanisms? Even though, there is not urgency on the study of CMB circular polarization, since the discovery of the CMB, there have been several attempts in the past and also at the present to experimentally measure it. Moreover, since CMB linear polarization has already been detected, the next step would be that of the study of circular polazation which as I will show in this paper is generated by very interesting mechanisms which are extremely important to the fundamental physics. The first studies on the CMB circular polarization were done in connection with studies on anisotropic expansion of the universe which are characterized by some type of Bianchi models \cite{negroponte:80}. Other studies on generation of CMB circular polarization include non commutative geometry \cite{Bavarsad:2009hm}, electron-positron scattering in magnetized plasma at decoupling time \cite{Giovannini:2010ar}, propagation of CMB photons in magnetic field of supernova remnants of the first stars \cite{De:2014qza}, photon-pseudoscalar mixing in magnetic field domains \cite{Agarwal:2008ac}, scattering of CMB photons with cosmic neutrino background \cite{Mohammadi:2013dea}. For a recent review on other CMB circular polarization mechanisms see Ref. \cite{King:2016exc}. The first experimental attempts to measure the circular polarization of the CMB were done in Ref. \cite{smooth:83} where no evidence for CMB circular polarization was found and only constraint on the degree of circular polarization was set. The current upper limit on the CMB circular polarization has been set by the MIPOL experiment \cite{Mainini:2013mja}, $P_C(T_0)\lesssim 7\times 10^{-5}-5\times 10^{-4}$ at the frequency 33 GHz and at angular scales between $8^\circ$ and $24^\circ$. In this work I study the impact of magneto-optic effects on the CMB polarization in the presence of cosmic magnetic fields where I mostly concentrate on generation of CMB circular polarization. A systematical study of the most important magneto-optic effects in the generation of a net CMB elliptic (circular and linear) polarization is done. By including all magneto-optic effects mentioned above, I derive the equations of motion for the Stoke's parameters which form a coupled system of differential equations. I use a density matrix approach to study the mixing of different magneto-optic effects and then solve the equations of motion by using perturbation theory. It turns out that among CM and vacuum polarization effects, the CM effect in plasma is the most promising effect in generation of elliptic polarization in the low frequency part of the CMB, while in the high frequency part, the vacuum polarization is the dominant one. I also will use current limit on the degree of circular polarization, to set new limits on the mass and coupling constant of pseudoscalar particles. In connection with CMB circular polarization, I calaculate its magnitude in terms of degree of circular polarization at present $P_C(T_0)$ and compare with experimental result(s). Generation and evolution of CMB E-mode and B-mode generated by the above mentioned effects is not studied in this work. This paper is organized as follows: In Sec. \ref{sec:2}, I derive the equations of motion for the photon and pseudoscalar fields in an expanding universe and introduce the photon polarization tensor in magnetized medium which describes forward scattering of photons. In Sec. \ref{sec:3}, I study the equations of motion for the density matrix in the case of open systems and establish the connection between the system Hamiltonian and the field mixing matrix. In Sec. \ref{sec:4}, I find the equations of motion for the density matrix in an expanding universe and solve them in the case of vacuum polarization and CM effects. In Sec. \ref{sec:5}, I present the equations of motion for the density matrix in the case when the contribution of the pseudoscalar field is included and introduce the concept of generalized Stokes parameters. Then I find perturbative solutions of the reduced Stokes vectors in transverse magnetic field. In Sec. \ref{sec:6}, I study the generation of CMB circular polarization in the case of photon-pseudoscalar particle mixing in transverse magnetic field and set new limits on the pseudoscalar parameter space. In Sec. \ref{sec:7}, I conclude. In this work I use the metric with signature $\eta_{\mu\nu}=\textrm{diag}(1, -1, -1, -1)$ and work with the natural (rationalized) Lorentz-Heaviside units ($k_B=\hbar=c=\varepsilon_0=\mu_0=1$) with $e^2=4\pi \alpha$.

\label{sec:7} In this work we have studied the most important magneto-optic effects and their impact in the generation of CMB polarization. We presented a systematic study of each of them where we mostly focused on the generation of CMB circular polarization. In this work we found the equations of motion for photon and pseudoscalar fields in an external magnetic field in the WKB approximation, and then found the equations of motion for the Stokes parameters by using density matrix approach as shown in Sec. \ref{sec:3}. The resulting equations describe the mixing of different magneto-optic effects which obviously complicate the situation but on the other hand give richer scenarios. In Sec. \ref{sec:4} we studied the vacuum polarization and CM effects separately, in order to isolate the contribution of each of them to CMB polarization. They are second order magneto-optic effects on magnetic field amplitude $B_e$ and are responsible for generation of phase shifts between the states $A_+$ and $A_\times$. These effects generate CMB elliptical polarization only in the case when the CMB is initially polarized. In this work we concentrated in the post decoupling epoch and worked under the hypothesis that the CMB acquired a small polarization at decoupling time due to Thomson scattering. We used perturbation theory and found the evolution as a function of $T$ of the Stokes parameters. We studied in particular the generation of circular polarization which is represented by the Stokes parameter $V$, in cases of observation angles $\Phi\neq\pi/2$ and $\Phi=\pi/2$. The contribution of vacuum polarization and CM effects to $V$ depends essentially on $\Phi$, $B_{e0}$, $\nu_0$ and on the magnitude of the Stokes parameters at decoupling which, on the other hand, depend on the temperature anisotropy. In this work we assumed that $V_i=0$ at decoupling while the other parameters are non zero. The magnitude of the parameters $Q_i$ and $U_i$ obviously are smaller than temperature anisotropy and observations of CMB linear polarization give an order of magnitude of $Q_i\sim U_i\sim 10^{-6}$. In the case of vacuum polarization and $\Phi\neq\pi/2$, the degree of circular polarization is proportional to $Q_i$ and $U_i$, as shown in \eqref{v-today} and in most cases is the term proportional to $U_i$ which dominates. This term on the other hand is proportional to $\nu_0$ and $B_{e0}^2$. Consequently, significant generation of circular polarization would occur in the high frequency part of the CMB and for higher values of $B_{e0}$. In this work we used in our estimates a canonical value of $B_{e0}\sim $ nG but in principle higher values are possible. If for example one observes the CMB in the Wien region, say at $\nu_0\simeq 700$ GHz and the magnetic field is of the order of 100 nG, the degree of circular polarization would be of the order $P_C\sim 10^{-11}$ while for $B_{e0}\sim $ nG is four orders of magnitude smaller. Also for the CM effect, the degree of circular polarization is proportional to the initial values of Stokes parameters at decoupling and to $B_{e0}, \nu_0$ and $\Phi$. One distinguishing feature of the CM effect is the relation between $V_0$ and $\nu_0$ which is $V_0\propto \nu_0^{-3}$. This relation makes the CM effect quite appealing in regard to generation of circular polarization in the Rayleigh-Jeans part of the spectrum. For $\Phi\neq\pi/2$ the degree of circular polarization is given in \eqref{CM-V} where the first term is proportional to $Q_i$ and the second term is proportional to $U_i$. The term proportional to $Q_i$ shares a common feature with the vacuum polarization by the fact it gets contribution from the Faraday effect which is encoded in $\rho$. Under the approximations used in Sec. \ref{sec:4}, the term proportional to $Q_i$ is smaller than that proportional to $U_i$. The latter coincides with the solution found for $V$ in case when $\Phi=\pi/2$ for $\mathcal G(T)\ll 1$ and $F(T)<1$. This means that the contribution of the CM effect to circular polarization is bigger in the limit $\Phi\rightarrow \pi/2$. The degree of circular polarization is substantive in the frequency region $\nu_0\sim 10^8-10^9$ Hz while for higher frequencies $\nu_0\sim 10^{11}$ Hz, the contribution of CM effect to $V_0$ is subdominant to vacuum polarization. For example, if $\nu_0\sim 10^{8}$ Hz and $B_{e0}\sim $ nG, the degree of circular polarization for the CM effect would be $P_C\sim 10^{-10}$ while if $B_{e0}\sim 100$ nG, $P_C\sim 10^{-6}$. In this work, we also studied the generation of elliptic polarization due to photon-pseudoscalar particle mixing in cosmic magnetic field, with emphasis on the degree of circular polarization. Differently from the vacuum polarization and CM effects, photon-pseudoscalar mixing has in addition two more independent parameters which are $m_\phi$ and $g_{\phi\gamma}$. We studied this mechanism in case of only transverse magnetic field and used perturbation theory to find the evolution in $T$ of the Stokes parameters. We used perturbation theory in two mixing regimes, namely weak and strong mixing and estimated the degree of circular polarization at present epoch. Since the parameters $g_{\phi\gamma}$ and $m_\phi$ are free and in general span a wide range of values, we used the present upper limit on the degree of circular polarization in order to constrain $g_{\phi\gamma}$ and $m_\phi$. These parameters on the other hand are constrained by the mixing regimes, therefore the limits that we presented are valid in these regimes. In the strong mixing regime, in general one has to solve trigonometric equations or inequations which have as independent variable $g_{\phi\gamma} B_{e0}$. The solutions generally, depend on an integer number $n$ and consequently they are not unique. The interval of values of $g_{\phi\gamma} B_{e0}$ can in principle be narrowed by complementary constraints on $g_{\phi\gamma}$ from other methods. On the other hand, in the weak mixing case there is not such a dependence on $n$. In this case by using the upper limit on $P_C$ obtained from MIPOL experiment, we got the constraint $|g_{\phi\gamma}|<4.29\times 10^{-19}(\textrm G/B_{e0})\quad \textrm{GeV}^{-1}$ for $m_\phi<1.6\times 10^{-14}$ eV. Other limits have been obtained late time generation of degree of linear polarization and by considering the case of almost unpolarized CMB at decoupling. In the weak mixing case, we obtained the average value over frequency of $\langle|g_{\phi\gamma}|\rangle\sim 10^{-18} ( \textrm{G}/B_{e0})$ for $m_\phi<1.6\times 10^{-14}$ eV and $P_L(T_0)\simeq 10^{-6}$. In the strong mixing case, again one obtains values of $g_{\phi\gamma} B_{e0}$ that depends on $n$ and therefore there is no unique solution. The same thing happens even in the resonant case with the particular case that, if, $g_{\phi\gamma} B_{e0}<1.17\times 10^{-24}$, then from $P_L\simeq 10^{-6}$ we get the value $|g_{\phi\gamma}|\simeq 1.66\times 10^{-27} ( \textrm{G}/B_{e0})$ for $m_\phi\simeq 1.6\times 10^{-14}$ eV. As in the case of vacuum polarization and CM effects, photon-pseudoscalar particle mixing generates non uniform degree of circular polarization and rotation of the polarization plane across the sky. This fact can be used in order to understand if the observed linear polarization at present has a non uniform component across the sky. If this would be true, it might be due to photon-pseudoscalar particle mixing if it is the dominant mechanism of generation of linear polarization among magneto-optic effects studied. From the experimental side, it turns out that among CM and QED effects, the CM effect is the most promising effect on generating circular polarization in the low frequency part of the CMB due to the dependence $V_0\propto \nu_0^{-3}$, while the vacuum polarization is the dominant one in the high frequency part due to $V_0\propto \nu_0$. The degree of circular polarization due to the CM effect in the low frequency part, in general, is bigger than that generated by vacuum polarization at high frequencies. Moreover, vacuum polarization and CM effects generate a rotation of the polarization plane of the CMB and this rotation together with the degree of circular polarization are not uniform across the sky because they depend on the observation angle $\Phi$. These facts would suggest that observation of CMB circular polarization is more likely to happen in the low frequency part of the CMB, mostly due to the CM effect and if, the observation frequency range is not a big detection issue. On the other hand, if one is interested in the measurement of the rotation angle of the polarization plane, the non uniformity of the rotation across the sky might be an issue. In order to detect CMB circular polarization, probably the most convenient frequency range would be for $\nu_0\sim 10^8-10^9$ Hz where the degree of circular polarization would be in the interval $P_C\simeq 10^{-13}-10^{-10}$ for $B_{e0}\simeq 1$ nG due to CM effect where the higher value of $P_C$ corresponds to the lower value of $\nu_0$. In this frequency range the contribution of vacuum polarization is completely negligible with respect to CM effect. The vacuum polarization is dominant to the CM effect in the Wien regime and for $\nu_0\sim 700$ GHz we found the interval $P_C\lesssim 10^{-15}-10^{-11}$ for the interval $B_{e0}\lesssim1-100$ nG. Assuming for the moment that frequency observation range is not an issue and it can be fixed based on experiment characteristic, one main problem on the detection of CMB circular polarization is related to the magnetic field amplitude which is poorly known. In the case of intergalactic magnetic field, usually upper limits are found by different methods and its amplitude is expected to be less than 1 nG (canonical value) up to a value of less than 100 nG. Using a canonical value of $B_{e0}\lesssim 1$ nG, one would expect that the degree of circular polarization to be $P_C\lesssim 10^{-10}$ in the Rayleigh-Jeans regime and $P_C\lesssim 10^{-15}$ in the Wien region. These values of the degree of circular polarization corresponds essentially to the case when $\Phi=\pi/2$ and for different values of $\Phi$, $P_C$ is usually smaller and not uniform across the sky. The photon-pseudoscalar mixing contributes to circular polarization as well and based on values of $g_{\phi\gamma}$ and $m_\phi$ found from current upper limit on circular polarization from MIPOL experiment, its contribution might be bigger than CM and vacuum polarization effects. So, let us assume for example that $B_{e0}\lesssim$ 1 nG and circular polarization would be detected with degree of circular polarization with value in the range, $10^{-10}\lesssim P_C\lesssim 10^{-6}$. This would mean there is a contribution to $P_C$ due to photon-pseudoscalar mixing which is much bigger than other magneto-optic effects or the amplitude of the magnetic field might be higher than assumed or circular polarization is generated by another effect not considered in this work. If one would detect CMB circular polarization with average value of $P_C\lesssim 10^{-10}$, it is more likely that CM and vacuum polarization effects are the source of this polarization. In all studied effects, we have assumed that the large scale magnetic field was present at the decoupling epoch therefore the field has been assumed to have primordial origin and a function of spacetime coordinates, $\bs B(\bs x, t)$, namely the field is non homogeneous in space and time. In the case of vacuum polarization, the expressions for the photon polarization tensor and derived quantities such as the index of refraction have been derived under the assumption that the electromagnetic field tensor satisfies the condition $|\partial_\mu F_{\sigma\rho}|\ll m_e |F_{\sigma\rho}|$, see Ref. \cite{Dunne:2004nc} for details. This condition on the electromagnetic field tensor translates into conditions on the magnetic field $|\partial B_e^i(\bs x, t)/\partial t|\ll m_e|B_e^i(\bs x, t)|$ and $|\partial B_e^i(\bs x, t)/\partial \bs x|\ll m_e|B_e^i(\bs x, t)|$. Obviously both the last conditions on $B_e^i(\bs x, t)$ are satisfied in an expanding universe for a large class of magnetic fields where the former can be written as $H^{-1}(t)\gg 2/m_e= 7.74\times 10^{-11}$ cm for the Hubble radius as a function of time and the latter condition can be written as $l_B(t)\gg 2/m_e= 7.74\times 10^{-11}$ cm where $l_B$ is the variation scale in space of the external magnetic field. In the case of the CM effect, the elements of photon polarization tensor are usually derived for constant magnetic fields but since this effect is similar to the QED effect, namely is of the second order in $B_e$, one can extend the results for constant fields also to the case when $|\partial_\mu F_{\sigma\rho}|\ll m_e |F_{\sigma\rho}|$ in complete analogy with the vacuum polarization effect. In the case of photon-pseudoscalar mixing, the magnetic field can be either homogeneous or non homogeneous as far as the photon wavelength $\lambda\ll l_B(t)$ in the WKB approximation, namely the magnetic field is a slowly varying function in space and time with respect to the photon wavelength or frequency. It is worth to mention also what has not been studied in this work. The first thing is related to Thomson scattering and scattering of pseudoscalar particles at post decoupling epoch, namely for $T<2970$ K and their absence in our density matrix formalism. In general, scattering is a mechanism of coherence breaking for mixing/oscillation processes which results in damping of the fields. In the density matrix formalism, the structure of the damping operator can be calculated by using field theory for scattering which is essentially the calculation of the commutator $[H_T, \rho]$ on the r. h. s. of \eqref{dens-eq} where $H_T$ includes the Hamiltonian for the Thomson scattering and that of scattering of pseudoscalar particles. However, quite often the damping term due to scattering, in case of non degenerate and non relativistic electron gas can be approximated\footnote{Similar situation occurs quite often in neutrino physics, see Ref. \cite{Dolgov:2002wy}.} by, $-i\{\Gamma, \rho\}$, where $\Gamma$ is the scattering rate matrix of photons and pseudoscalar particles which is diagonal in the basis $|A_+\rangle, |A_\times\rangle, |\phi\rangle$. Consequently, the damping term due to scattering, would have the same structure as the damping term due to Hubble friction. Therefore, the Stokes parameters would be affected by scattering but not their ratio because it cancels out exactly as the damping term due to Hubble friction. However, one must always keep in mind that this is an approximation. The second thing is related to the case $\Phi\neq\pi/2$ for the photon-pseudoscalar particle mixing. In Sec. \ref{sec:5}, we found the equations of motion for the reduced Stokes vectors in case of transverse external magnetic field. In this case it was possible to find two sets of decoupled differential equations for the reduced Stokes vectors and solved the equations by using perturbation theory. If the field is not transverse, namely $\Phi\neq\pi/2$, in general one has to solve simultaneously, a system of nine linear differential equations of the first order which can be problematic to solve even numerically because quite often they are stiff. We shall treat this problem in more details elsewhere but even at this stage we can outline very important conclusions about the nature of the solutions and the impact on the CMB polarization. The importance of solutions of the equations of motion in the case $\Phi\neq\pi/2$ (for photon-pseudoscalar mixing) relies in the fact, that being the system of equations linear, see Eqs. \eqref{fin-system}, the solutions will be proportional to initial values at a given temperature $T_i$ which does necessarily coincides with decoupling temperature. Therefore, each Stokes parameter would be proportional to $I_\gamma(T_i), Q_i, U_i, V_i, S_{4i}$ etc., and for $T<T_i$ usual Stokes parameters (those which in general interest us) would be different from zero even in case of initially unpolarized CMB at $T=T_i$. We saw similar situation in Sec. \ref{sec:6.3}, where we studied the case of unpolarized CMB at decoupling for transverse magnetic field. Consequently, the CMB would acquire polarization independently on Thomson scattering, even in case when it is initially unpolarized. This situation would be very important in order to investigate prior decoupling CMB polarization due to photon-pseudoscalar mixing in external magnetic field. According to standard cosmology, generation of CMB polarization occurs at or very close to decoupling time due to Thomson scattering when the condition of tight coupling between photons and electro-baryon plasma breaks down. Indeed, for most models of generation of CMB polarization which include scalar perturbations, magnetic fields, gravitational waves etc., at the end is always the Thomson scattering which generates CMB polarization \cite{Zaldarriaga:2003bb}. The tight coupling condition would imply that, if there is any degree of polarization prior to decoupling, generated at temperature $T$, it would be damped very fast due to scattering of photons with electrons and baryons. However, as we have seen and discussed in this work, photon-pseudoscalar mixing apart from generating temperature anisotropy as shown in Sec. \ref{sec:6.3} and spectral distortions of the CMB \cite{Mirizzi:2009nq}, it generates also polarization, independently on Thomson scattering. Consequently, here we advance the hypothesis that photon-pseudoscalar mixing \emph{might} generate non uniform CMB polarization across the sky, even before decoupling epoch, if the rate of photon-pseudoscalar oscillation is faster than photon scattering rate with electro-baryon plasma. Obviously, all said about this hypothesis would depend on pseudoscalar field parameters $m_\phi$ and $g_{\phi\gamma}$. The suggested hypothesis needs further attentive study and it would be too premature to conclude that it is indeed the case. \vspace{1cm} {\bf{AKNOWLEDGMENTS}}: This work is supported by the Russian Science Foundation Grant Nr. 16-12-10037. I would like to thank LNGS for the support received through the fellowship POR 2007-2013 `Sapere e Crescita' where part of this research was conducted. \appendix

1607.02094

1607

1607.07723_arXiv.txt

{The evolution of the universe during the Dark Ages (DA) and the Epoch of Reonization (EoR) marks an important transition in the history of the universe but it is not yet fully understood.} {We study here an alternative technique to probe the DA and EoR that makes use of the Comptonization of the CMB spectrum modified by physical effects occurring during this epoch related to the emergence of the 21-cm radiation background. Inverse Compton scattering of 21-cm photon background by thermal and non-thermal electrons residing in the atmospheres of cosmic structures like galaxy clusters, radiogalaxy lobes and galaxy halos, produces a specific form of Sunyaev-Zel'dovich effect (SZE) that we refer to as SZE-21cm.} {We derive the SZE-21cm in a general relativistic approach which is required to describe the correct spectral features of this astrophysical effect. We calculate the spectral features of the thermal and non-thermal SZE-21cm in galaxy clusters and in radiogalaxy lobes, and their dependence on the history of physical mechanisms occurring during the DA and EoR. We study how the spectral shape of the SZE-21cm can be used to establish the global features in the mean 21-cm spectrum generated during and prior to the EoR, and how it depends on the properties of the (thermal and non-thermal) plasma in cosmic structures. } {We find that the thermal and non-thermal SZE-21cm have peculiar spectral shapes that allow to investigate the physics and history of the EoR and DA. Its spectrum depends on the gas temperature (for the thermal SZE-21cm) and on the electrons minimum momentum (for the non-thermal SZE-21cm). The global SZE-21cm signal can be detected (in $\sim 1000$ hrs) by SKA1-low in the frequency range $\nu \simgt 75-90$ MHz, for clusters in the temperature range 5 to 20 keV, and the difference between the SZE-21cm and the standard SZE can be detected by SKA1 or SKA2 at frequencies depending on the background model and the cluster temperature. } {We have shown that the detection of the SZE-21cm can provide unique information on the DA and EoR, and on the cosmic structures that produce the scattering; the frequencies at which the SZE-21cm shows its main spectral features will indicate the epoch at which the physical processes related to the cosmological 21-cm signal occurred and shed light on the cosmic history during the DA and EoR by using local, well-known cosmic structures like galaxy clusters and radio galaxies.}

Departures of the Cosmic Microwave Background (CMB) frequency spectrum from a pure blackbody encode information about the thermal history of the early universe before the epoch of recombination when it emerged from the last scattering surface. The evolution of the universe after this epoch proceeds through the period of the Dark Ages (DA) that ends $\sim 400$ million years later, when the first galaxies formed and start emitting ionizing radiation. The transition period at the end of the DA marks the Epoch of Reionization (EoR). During this epoch, radiation from the very first luminous sources (e.g., early stars, galaxies, and quasars) succeeded in ionizing the neutral hydrogen gas that had filled the universe since the recombination event (see, e.g., Barkana and Loeb 2001, Loeb and Barkana 2001, Bromm and Larson 2004, Ciardi and Ferrara 2005, Choudhury and Ferrara 2006, Furlanetto et al. 2006, Morales and Wyithe 2010). The current constraints suggest that the EoR roughly occurs within the redshift range of $z \approx 6 - 20$. This cosmic period is not yet completely understood and various astrophysical probes have been suggested to shed light on this epoch for early structure formation (see Zaroubi 2013 for a review). Information from the DA period is not explicitly contained in the CMB because baryonic matter and radiation have already decoupled, and the bulk of baryonic matter in the universe during this period is in the form of neutral hydrogen gas in the inter galactic medium (IGM). Rather than target observations at the first galaxies and quasars that are the rare, early products of gravitational collapse, it is then necessary to detect directly the presence of the ubiquitous hydrogen gas. One of the methods of achieving this detection is to search for signatures of the (highly redshifted) 21-cm hyperfine transition line of neutral hydrogen (see, e.g., Loeb \& Zaldarriaga 2004; Cooray 2004; Bharadwaj \& Ali 2004; Carilli et al. 2004; Furlanetto \& Briggs 2004; Furlanetto et al. 2006; Pritchard \& Loeb 2010, 2012; Liu et al. 2013). The 21-cm signal from the DA would appear as a faint, diffuse background detectable at frequencies below 200 MHz (for redshifts $z > 6$). Thus, measuring the brightness temperature of the redshifted 21-cm background could yield information about both the global and local properties of the IGM. Determining the average brightness temperature over a large solid angle as a function of redshift would eliminate any dependence on local density perturbations and constrain the history of the neutral fraction of hydrogen in the IGM. It has been noted that there are several problems related to the observation of the 21-cm background. Firstly, this signal is faint, of the order of tens of mK relative to the CMB (see, e.g., Furlanetto et al. 2006), and until now only upper limits have been obtained (see, e.g., Paciga et al. 2013, Dillon et al. 2014, Parsons et al. 2014). The second problem is related to the presence of galactic and extragalactic foregrounds whose amplitude can be also about four order of magnitude larger than this signal (see, e.g., de Oliveira-Costa et al. 2008). These problems make difficult to study this signal with the present-day and new generation of radio interferometers, since they are not sensitive to the mean signal, but only to its inhomogeneity, and thus require a very precise calibration and knowledge of foregrounds to remove their contribution (see, e.g., discussion in Furlanetto et al. 2006). Various methods have been proposed to overcome these problems. One possibility is to study the 21-cm fluctuations to measure the mean background through their redshift-space anisotropies (Barkana \& Loeb 2005a); this method can be used with the next generation instruments like the Square Kilometer Array (SKA) (see, e.g., McQuinn et al. 2006). A second method is to measure the contrast between the 21-cm signal and the bubbles of ionized plasma present during the EoR, and use their contrast to measure the mean amount of neutral gas (see, e.g., Furlanetto et al. 2006 and references therein). An alternative method that we want to discuss extensively in this paper is to use the SZE-21cm, i.e. the spectral distortion of the CMB spectrum modified by physical effects occurring during the epoch related to the emergence of the 21-cm radiation background, induced by inverse Compton scattering off the intervening electrons in the atmospheres of various cosmic structures, like galaxy clusters, radiogalaxy lobes and galactic halos. A preliminary attempt to calculate the SZE-21cm has been presented by Cooray (2006). This calculation turns out to be not adequate for a correct description of the SZE-21cm for two reasons:\\ i) the photon background model used for the modification to the CMB caused by mechanisms working during the DA and EoR is unphysical, because it contains a number of artificial discontinuities, under-resolves the main features of interest at $\nu\sim70$ MHz and contains an unphysical reionization history that produces substantial 21-cm signal down to redshifts $z<2$ (i.e., at frequencies $> 300$ MHz);\\ ii) it is performed in the non-relativistic approximation of the Compton scattering process of CMB photons in the hot intra-cluster medium of galaxy clusters thus neglecting any effect induced by the relativistic corrections to this scattering, by multiple scattering effects and by the scattering of additional non-thermal electrons in clusters, as explicitly reported by Cooray (2006).\\ Such problems in the Cooray (2006) calculations lead to an incorrect description of the SZE-21cm that has important consequences in using this cosmological probe. In fact, to take full advantage of the SZE-21cm study, it is necessary to use a full relativistic formalism, its generalization to any order of magnitude in the plasma optical depth $\tau$ and the possibility to include also the combination of various electron populations residing in cosmic structures (see, e.g., Colafrancesco et al. 2003). It is also necessary to use a wider and more physically motivated set of models for the 21-cm background, including also other physical processes that can change this background, such as the effect of Dark Matter heating. Finally, it is worth considering the effect of changing the redshifts at which the different physical processes took place. In this paper we perform such a more complete study following the previous lines of investigation.\\ First, to describe the CMB spectrum modified by the 21-cm cosmological background, we use the results of the 21cmFAST code (Mesinger, Furlanetto \& Cen 2011) that include realistic physical effects and also additional mechanisms, such as the heating induced by Dark Matter annihilation (e.g., Valdes et al. 2013; Evoli et al. 2014).\\ Secondly, we perform the calculations in the full relativistic formalism for the derivation of the SZE (see, e.g., Colafrancesco et al. 2003 for details), that is suitable to calculate the SZE-21cm in detail, and to derive the precise information about its spectral properties over a wide frequency range and in a wide set of cosmic structures. This general treatment allows, therefore, to increase both the number and the redshift distribution of objects that can be studied with this method, including galaxy clusters with high temperatures (which are the best targets for maximizing the SZE-21cm signal and are more subject to relativistic effects), with radio halos, cool-cores and other complex morphologies, as well as other extragalactic sources with non-thermal electron distributions such as radio galaxies lobes. The plan of the paper is the following: in Sect 2 we present the general, full relativistic derivation of the SZE-21cm and the models for the frequency distribution of the global 21-cm background we use in the paper. These are new crucial elements of the derivation of the SZE-21cm that have never been provided up to date. In Sect. 3 we discuss the results of our calculations for various scenarios of the radiation background emerging from the DA and EoR, considering various astrophysically motivated scenarios. We also discuss here, for the first time, the derivation and the possibility to observe both the thermal and the non-thermal SZE-21cm. We discuss our results in the light of the future radio interferometric experiments like the SKA in Sect.4, and we summarize our conclusions in Sect.5. Throughout the paper, we use a flat, vacuum--dominated cosmological model with $\Omega_m = 0.315$, $\Omega_{\Lambda} = 0.685$ and $H_0 =67.3$ km s$^{-1}$ Mpc$^{-1}$.

The goal of obtaining information on the physical processes occurred during the DA and EoR by measuring the SZE-21 cm with SKA is challenging, but possible if pursued with good theoretical and observational strategies. Observations have to be carried out towards high temperature and high optical depth clusters to maximize both the overall signal and the difference between the standard and the modified SZE. The best frequency ranges of observation of the SZE-21cm are between $\sim$ 90 and 120 MHz, where the difference between the standard and the modified SZE is maximum. In our benchmark model, the sensitivity of SKA1-low is good enough to detect this difference with 1000 hours of integration, whereas for the other background models the difference between the standard and the modified SZE can be detected only with SKA2 for the same integration time in frequencies bands that depend on the background model and the temperature of the cluster. Together with very deep observations, a very accurate theoretical analysis is required, where the full formalism to calculate the SZE and detailed models for describing the effect of the cosmological 21-cm background on the CMB spectrum have to be used. In addition, we find that a very important strategy will be the detailed study of the SZE at higher frequencies in order to estimate the gas parameters to be used as prior constraints for the study of the SZE-21 cm at low frequencies.\\ Observations in the frequency bands of SKA1-mid are also very important to disentangle the SZE from the cluster synchrotron emission. In this respect, the use of high-redshift clusters can alleviate the problem, since the radio emission decreases as $D_L^{-2}$, whereas the SZE is not depending on the cluster distance. The detection of the non-thermal SZE-21 cm appears to be more challenging, since the signal is much fainter with respect to the thermal one, especially regarding the difference between the standard and the modified SZE, that can be also a factor of $\sim 10^2$ smaller with respect to the thermal case. However, the different spectral features can allow, in principle, a detection of this signal and hence an estimate of non-thermal cluster properties independently of measurements in other spectral bands. We note here that it is possible to strategize the search of this signal in objects where the non-thermal components are dominant, such as in the case of radio galaxies lobes. In this case, objects with more energetic electrons (i.e. with harder radio spectra), large optical depth (for which a good indication could be a strong radio luminosity) and high redshift are preferable. The independence of the SZE from the redshift can allow the study of the SZE-21cm in a larger number of objects spread over a wider redshift range, therefore producing statistical studies aimed at maximizing the detectable signal, and detect the properties of the 21-cm background and of the early DM halos over a large set of spatial directions, allowing in such a way a better understanding of the full cosmic history of the physical processes occurring in the Dark Ages and the Epoch of Reionization.

1607.07723

1607

1607.05383_arXiv.txt

We present the first scientific results from the luminous red galaxy sample (LRG) of the extended Baryon Oscillation Spectroscopic Survey (eBOSS). We measure the small and intermediate scale clustering from a sample of more than 61,000 galaxies in the redshift range $0.6 < z < 0.9$. We interpret these measurements in the framework of the Halo Occupation Distribution. The bias of eBOSS LRGs is $2.30 \pm 0.03$, with a satellite fraction of $13\pm3$\% and a mean halo mass of $2.5\times10^{13}h^{-1}M_{\odot}$. These results are consistent with expectations, demonstrating that eBOSS galaxies will be reliable tracers of large scale structure at $z\sim 0.7$. The eBOSS galaxy bias implies a scatter of luminosity at fixed halo mass, $\sigma_{\log L}$, of 0.19 dex. Using the clustering of massive galaxies from BOSS-CMASS, BOSS-LOWZ, and SDSS, we find that $\sigma_{\log L}=0.19$ is consistent with observations over the full redshift range that these samples cover. The addition of eBOSS to previous surveys allows investigation of the evolution of massive galaxies over the past $\sim 7$ Gyr.

Galaxy redshift surveys have been fundamental in advancing our understanding of the universe. The successes of the past decade, varying from 2dFGRS \citep{Cole_2005}, SDSS \citep{Eisenstein_2005, Zehavi_2011}, and BOSS \citep{Anderson_2012}, have spawned even larger investments in mapping the universe through the three-dimensional distributions of galaxies. In this paper, we present the first measurements of the clustering of luminous red galaxies (LRGs) from the extended Baryon Oscillation Spectroscopic Survey (eBOSS; \citealt{eBOSS_Dawson}), the successor program to BOSS (\citealt{Dawson_BOSS}). The eBOSS LRG program has the power to provide reliable measurements of galaxy clustering. We focus on LRG clustering at small scales ($r \lesssim 20$ $h^{-1}$ Mpc), scales which provide information on the bias of the galaxy sample and how these galaxies are distributed in dark matter halos. The framework in which we interpret the eBOSS data is the Halo Occupation Distribution (HOD). This approach describes the bias relation between the galaxies and matter at the level of ``virialized" dark matter halos which are expected to be in approximate dynamical equilibrium \citep{HOD_Weinberg, Peacock_2000, Seljak_2000, Benson_2000, Martin_2001, Cooray_2002}. In the HOD framework, the key quantity is the probability distribution $P(N|M)$ that a halo of virial mass $M$ contains $N$ galaxies of a given type, along with the relations between the galaxy and dark matter spatial and velocity distributions within halos. Given an HOD and a particular cosmological model, the statistics of galaxy clustering can be predicted in the sense that the cosmological model determines the properties of the halo distribution, while the HOD specifies how those halos are populated with galaxies. HOD modeling has been used to interpret clustering in nearly all large-scale galaxy redshift surveys (e.g. \citealt{Zheng_DEEP2, Zheng_2009, Zehavi_2011, CMASS_Martin, Parejko_LOWZ, Guo_2014}). The HOD results provide physically informative and important information to test theories of galaxy formation and evolution. One of the key quantities in galaxy formation is the scatter in galaxy luminosity (or stellar mass) at fixed halo mass. Clustering is one of the few methods that is sensitive to the scatter. We will use the HOD to estimate this scatter and compare it to other galaxy samples spanning a redshift range of $z=0.7$ to $z=0.1$. We will show that this scatter is both small (0.19 dex in $\log{L}$) and constant over this redshift range. Our paper is organized as follows. Section 2 briefly describes the eBOSS observations and the definition of our LRG sample. The measurement of clustering is presented in Section 3, along with the comparison with the BOSS result. In Section 4, we interpret our result in the framework of HOD. Finally, the conclusion and discussion of our measurements as well as its implication are given in Section 5. Throughout this paper, the distances are measured in units of $h^{-1}$ Mpc with the Hubble constant $H_{0}$=100 $h$ km s$^{-1}$ Mpc$^{-1}$. The redshifts are converted to distances by assuming a spatially flat $\Lambda$CDM model with $(\Omega_{m}, h)$ = $(0.29, 0.7)$. The same cosmology is also used for the $N$-body simulations to make mock catalogs. The halos are defined as the spherical overdensity masses which are 200 times the background density.

This paper marks the first scientific results from the eBOSS LRG program. Although the observing strategy for eBOSS LRGs differs substantially from its predecessors in BOSS and SDSS, we have demonstrated that the combination of the bright-end of the BOSS CMASS sample with the eBOSS LRGs over the redshift range $0.6<z<0.9$ provides a robust clustering sample at small and intermediate scales. Our halo occupation analysis of this sample indicates that these galaxies have properties that are well-placed within our understanding of the relationship between massive galaxies and dark matter halos, with a bias factor of $b=2.30$, a satellite fraction of $\sim 13\%$, and halo mass scale in agreement with the scaling relations calibrated on other surveys. The addition of the eBOSS galaxy sample to previous spectroscopic samples yields a set of massive galaxies that span that last $\sim 7$ Gyr of the history of the universe. Our measurement of scatter in galaxy luminosity at fixed halo mass, $\sigma_{\log{L}}=0.19\pm0.02$, is in good agreement with other studies that have focused on $z=0$ samples. \cite{lehmann_etal:15}, using galaxy clustering alone, reported a value of $0.17^{+0.03}_{-0.05}$; \cite{reddick_etal:13}, using a combination of galaxy groups and clustering, find $0.21^{+0.01}_{-0.02}$; and \cite{more_etal:09}, using satellite kinematics, find $0.16\pm0.04$. Assuming these measurements are all independent (which is not strictly true), the weighted combination of all four results indicate $\sigma_{\log{L}}=0.19\pm 0.01$, a value that is somewhat larger than recent measurements of the scatter in stellar mass at fixed halo mass, $\sigma_{\log{M\ast}}\approx 0.16$ (\citealt{li_etal:12, kravtsov_etal:14, tinker_etal:16_boss, zu_mandelbaum:16}), which itself appears to be independent of redshift. The larger scatter in luminosity, for galaxies that are nearly all on the red sequence, is indicative of different formation histories at fixed stellar mass that yield different stellar-$M/L$ ratios and mean stellar ages. At first glance, the lack of evolution of either scatter value is notable but not surprising given that the massive end of the red sequence is constructed prior to $z\sim 1$ and that massive galaxies evolve in a manner close to passive stellar evolution over that timespan (\citealt{wake_etal:08, cool_etal:08}). However, true passive evolution of massive galaxies would result in a reduction in $\sigma_{\log{L}}$ as galaxies evolve, due to the fact that $M/L$ ratios for passive stellar populations evolve to the same asymptotic value. To match dynamically passive evolution, $\sigma_{\log{L}}$ would have to decrease from 0.19 at $z=0.7$ to 0.12 at $z=0.1$, which is clearly ruled out by our measurements. \cite{Gu_2016} find that the scatter (in stellar mass) induced by hierarchical merging is constant with redshift, but merging is not the dominant source of scatter at the halo masses probed by eBOSS galaxies. For galaxies in $10^{13}$ $h^{-1}M_{\odot}$ halos, in-situ star formation is still predicted to be the dominant source of scatter. Abundance-matching studies by \cite{Behroozi_2013b} and \cite{moster_etal:13} demenstrate that stellar mass growth from merging accounts for $\sim 10\%$ of the $z=0$ galaxy mass. This result is in agreement with earlier clustering studies of massive galaxies that found LRG merger rates of $\sim 1\%$ per Gyr (\citealt{wake_etal:08} and references therein). How does a population without merging or star formation have a constant luminosity scatter for over half the lifetime of the universe? SDSS, CMASS, LOWZ, and eBOSS represent a heterogeneous set of galaxy samples. Our SDSS sample is volume-limited, and at $M_r<-21.7$ the fraction of star-forming objects is negligible. The BOSS samples, as a whole, suffer from high significant incompleteness due to their color-based selections (\citealt{leauthaud_etal:16, tinker_etal:16_boss}), but by using only the brightest third of each sample in relatively narrow redshift ranges, CMASS and LOWZ are roughly complete as well. eBOSS, however, cannot be considered a complete sample. It is not trivial to estimate what the bias of a complete eBOSS sample would be at the number density used to create our subsamples, $1.4\times 10^{-4}(h^{-1} \text{Mpc})^{-3}$. The color selection excludes some brighter galaxies and includes some fainter objects, but the fainter objects will be redder and thus possibly more clustered than the brighter, but bluer, excluded objects. This is true of the overall CMASS sample (c.f. Figure 7 of \citealt{tinker_etal:16_boss}). If this is true of eBOSS, then the overall trend of $b(z)$ in Figure \ref{bias} would be consistent with some small reduction in $\sigma_{\log L}$ with time. Alternatively, the scatter in stellar $M/L$-ratio on the red sequence may not change enough between $z=0.7$ and $z=0.1$ to be detectable within our precision of 0.02 dex in scatter, since this scatter would add in quadrature with the scatter in stellar mass at fixed halo mass. Stellar population synthesis models would be required to address this question within the precision of our measurements, and will be included in a future work. The primary science driver of the eBOSS LRG sample is to probe the growth and expansion history of the universe at $z=0.7$. As a part of the SDSS-IV project, the eBOSS survey takes over the mission from its precursor BOSS and will map the universe in a higher redshift range and larger volume. After roughly one year observation, we reach a LRG sample with more than 34000 massive galaxies at an effective redshift $z\sim 0.7$. The result here shows that eBOSS is working well and the designed expectation is being reached. The clustering measurements that will be achieved with this sample through the completion of this survey will an important extension toward a complete map of the observable universe.

1607.05383

1607

1607.00874_arXiv.txt

The relations between observable stellar parameters are usually assumed to be deterministic. That is, given an infinitely precise measurement of independent variable, `$x$', and some model, the value of dependent variable, `$y$' can be known exactly. In practice this assumption is rarely valid and intrinsic stochasticity means that two stars with exactly the same `$x$', will have slightly different `$y$'s. The relation between short-timescale brightness fluctuations (flicker) of stars and both surface gravity \citep{bastien:2013} and stellar density \citep{kipping:2014} are two such stochastic relations that have, until now, been treated as deterministic ones. We recalibrate these relations in a probabilistic framework, using Hierarchical Bayesian Modelling (HBM) to constrain the intrinsic scatter in the relations. We find evidence for additional scatter in the relationships, that cannot be accounted for by the observational uncertainties alone. The scatter in surface gravity and stellar density does not depend on flicker, suggesting that using flicker as a proxy for $\log g$ and $\rho_\star$ is equally valid for dwarf and giant stars, despite the fact that the observational uncertainties tend to be larger for dwarfs.

\label{sec:intro} Accurate stellar characterization plays a vital role for many active research fields within astronomy. For example, stellar populations, galactic archaeology, the study of binary stars, asteroseismology and exoplanet studies all rely on inferences of basic stellar parameters to varying degrees. Empirically-derived and reliable estimates are of particular value, increasing our confidence in the end-product results built upon these inputs. Basic stellar parameters, such as effective temperature and surface gravity, can be inferred using one (or more) of several types of observations, such as spectroscopy, photometry, interferometry, etc. This inference can be performed by invoking theoretical models or by building an empirical calibration library. For example, an observed stellar spectrum could be matched against a library of theoretical spectra generated using stellar atmosphere models, or, against a library of observed spectra of ``standard stars'', serving as calibrators. Regardless of the approach, be it theoretical or empirical, the methods used for the inference of stellar parameters are traditionally ``deterministic''. In this context, a deterministic model can be loosely described as one where a particular observational input always returns a single-valued output for a parameter of interest, i.e. nature itself has no variance and the underlying model is considered to be a perfect description of reality. An alternative approach for inferring model parameters is to allow relationships between observables to be stochastic. In recent years, there has been a shift towards such methods in several areas of astronomy, particularly within the exoplanet community. For example, \citet{wolfgang:2015} considered that the mass-radius relationship of exoplanets is stochastic, since a particular sized planet could be have a range of planet masses due to unmodeled variances in compositions, environment and other complications. These recent demonstrations in exoplanetary science have prompted us to consider the need for treating the parent stars in the same probabilistic framework, with potential applications spanning many fields of astronomy. The demand for probabilistic stellar parameters is not only motivated by the fact that probability distributions are far more representative of our `beliefs' about astrophysical parameters, it also has a practical purpose. When using data published in the astronomical literature to, for example, infer relationships between parameters that are themselves the product of an inference process (for example, exoplanet transit depth and period), inference can be performed as the final stage in a hierarchical treatment \citep[see, e.g.][]{foreman-mackey:2014}. Studies such as these are benefited by posterior PDF samples, rather than point estimates of inferred properties. One of the more recent tools developed to characterize stars is known as ``flicker'' \citep{bastien:2013}. Flicker is a proxy for the scatter on an 8-hour timescale (denoted as $F_8$) in a broad visible bandpass time series photometric light curve, such as that from \textit{Kepler} or the upcoming TESS mission. A more detailed account of the proceedure to calculate flicker is described in \citet{bastien:2013}. As shown in \citet{bastien:2013}, flicker displays a remarkable correlation to the asteroseismically determined parent star surface gravities (\logg). Turning this around, the observation implies that flicker can be used to empirically infer surface gravities at the level of $\sim0.1$\,dex, an attractive proposition given the wealth of photometric light curves available through the array of exoplanet transit missions flying and scheduled to launch. \citet{cranmer:2014} demonstrated that models of stellar surface granulation indeed reproduce a flicker effect in close agreement with that observed by \citet{bastien:2013}, providing a physically-plausible explanation. Since surface gravity is highly correlated with mean stellar density (\rhostar) on evolutionary tracks, \citet{kipping:2014} showed that flicker can be also be used to infer \rhostar, which is more useful for exoplanet transit analysis \citep{seager:2003}. Whether one calibrates flicker to \logg\ or \rhostar, there are several aspects of the problem which are attractive for our purposes of a simple demonstration of probabilistic inference of stellar parameters. Firstly, in log-log space the relationship is very simple, appearing to be linear \citep{kipping:2014}. Secondly, there is a sufficiently large number of points in the sample (439 stars) to constrain a population-based model. Thirdly, there is significant excess scatter around the best-fitting relation implying that a deterministic model is inadequate. This is not surprising given that granulation is a complex and messy process for which one should not expect any parametric model to provide a perfect description. Finally, the physical processes that produce surface granulation, of which flicker is an observational tracer, may be more or less noisy for different types of stars. We will test whether flicker has greater predictive power in certain regions of parameter space; i.e. is flicker significantly more informative for subgiants than for dwarfs? For these reasons, we identify the calibration of flicker to \logg\ and \rhostar\ as a well-posed problem to first demonstrate probabilistic inference in the arena of stellar characterization.

\label{sec:discussion} We have recalibrated the relation between short timescale brightness fluctuations in the {\it Kepler} light curves of stars (flicker) with both stellar density and surface gravity, whilst including parameters to describe the intrinsic scatter in these relationships, presented in table \ref{tab:results}. The terms $\sigma_\rho$ and $ \sigma_g$ are both non-zero, suggesting that there {\it is} an additional source of scatter in the relations, not accounted for by the observational uncertainties alone. This is either caused by intrinsic scatter in the physical relationship between flicker and density and \logg, produced by some physical process that is not accounted for in the model, or by an underestimation of the observational uncertainties. We also tested a model with both an additional variance term {\it and} a term that included flicker-dependent variance. We found that the need for additional flicker-dependent variance was not supported by the data, indicating that the intrinsic scatter in the relations between flicker, \logg\ and \rhostar\ does not depend on evolutionary state. This is a simple `fitting a line to data' exercise, however it continues the discussion of probabilistic modelling that is an active topic within the fields of exoplanet and stellar astronomy. We used Hierarchical Bayesian Modelling (HBM) to constrain the intrinsic scatter in the relationship between flicker, surface gravity and density and included the effects of the non-negligible two-dimensional observational uncertainties. Relationships between astronomical parameters are almost always non-deterministic; an element of stochasticity effects the physical parameters of stars so one can never perfectly predict $y$ given an observation of $x$. We advocate a probabilistic approach in both the `fitting the model to data' step, {\it and} when using an empirically calibrated model to predict parameter values. The fitting stage benefits because if the relationships between parameters are falsely assumed to be deterministic, they will be skewed by data points with uncertainties that only represent measurement error and no additional scatter. The prediction stage benefits from the stochastic treatment both because a probability distribution is in many ways more representative of an observation than a point estimate, and because posterior PDF samples can be used in subsequent studies (provided the prior used during the fitting process is described). We provide posterior PDF samples at \url{https://zenodo.org/deposit/105051/}. Whenever a prediction for the surface gravity or density of a star is required, for a given estimate of flicker, we recommend using these posterior samples within the calculation of \rhostar or \logg\ and its (Monte Carlo) uncertainty. These posterior samples will naturally fold in the covariances between parameters. Simple analytical uncertainty propagation is only valid when uncertainties are Gaussian and uncorrelated which is rarely true and certainly not the case when the model is a straight line (the slope and intercept are alway correlated). A flicker value with uncertainties (or even better: posterior PDF samples), input into our model will result in a probability distribution over stellar densities or surface gravities which reflects both the uncertainties on the flicker measurement, the uncertainties on the model parameters {\it and} the intrinsic scatter in the flicker-\rhostar-\logg\ relations.

1607.00874

1607

1607.08162_arXiv.txt

During the planet formation process, billions of comets are created and ejected into interstellar space. The detection and characterization of such interstellar comets (also known as extra-solar planetesimals or extra-solar comets) would give us \emph{in situ} information about the efficiency and properties of planet formation throughout the galaxy. However, no interstellar comets have ever been detected, despite the fact that their hyperbolic orbits would make them readily identifiable as unrelated to the solar system. Moro-Mart{\'{\i}}n et al. 2009 have made a detailed and reasonable estimate of the properties of the interstellar comet population. We extend their estimates of detectability with a numerical model that allows us to consider ``close'' interstellar comets, e.g., those that come within the orbit of Jupiter. We include several constraints on a ``detectable'' object that allow for realistic estimates of the frequency of detections expected from the Large Synoptic Survey Telescope (LSST) and other surveys. The influence of several of the assumed model parameters on the frequency of detections is explored in detail. Based on the expectation from Moro-Mart{\'{\i}}n et al. 2009, we expect that LSST will detect 0.001-10 interstellar comets during its nominal 10-year lifetime, with most of the uncertainty from the unknown number density of small (nuclei of $\sim$0.1-1 km) interstellar comets. Both asteroid and comet cases are considered, where the latter includes various empirical prescriptions of brightening. Using simulated LSST-like astrometric data, we study the problem of orbit determination for these bodies, finding that LSST could identify their orbits as hyperbolic and determine an ephemeris sufficiently accurate for follow-up in about 4--7 days. We give the hyperbolic orbital parameters of the most detectable interstellar comets. Taking the results into consideration, we give recommendations to future searches for interstellar comets.

The current understanding of planet formation suggests that very large numbers of minor bodies are ejected into interstellar space by planets during and after formation \citep[e.g.,][]{1972IAUS...45..329S,1987AJ.....94.1330D,2004come.book..153D}. In typical simulations of solar system formation, only a small fraction of the small bodies that have a close encounter with the giant planets are captured into the Oort cloud \citep[the current source of long-period comets,][]{Oort1950}, the rest are ejected into the interstellar medium. Despite changing understanding of the formation and properties of our Oort cloud \citep[e.g.,][]{Levison2010,2011Icar..215..491K}, extra-solar debris disks \cite[e.g.,][]{2012Natur.490...74D}, and planet formation \citep[e.g.,][]{2012arXiv1211.1673C}, there is general consensus that interstellar space must be populated with a non-trivial population of small bodies, including those corresponding in size to the asteroids and comets in the solar system. Minor planets that originate in other planetary systems but are currently unbound are usually called interstellar comets (ICs)\index{interstellar comet}.\footnote{Although there are no codified definitions, objects unbound from any star that are the natural minor body extension of the interstellar medium are usually called interstellar comets \citep[e.g.,][]{1975AJ.....80..525W,1976Icar...27..123S}, while minor bodies detected orbiting around other stars are a natural extension of extra-solar planets and are often referred to as extra-solar planetesimals\citep[e.g.,][]{2005AJ....130.1261J}. \citet{M09} is an exception to this typical nomenclature.} Although there is one candidate IC known ($\S$\ref{sec:candidate}), at present, their existence is essentially theoretical \citep{1975AJ.....80..525W,1976Icar...27..123S, McGlynn89, Francis05}. Technically, many long period comets have slightly hyperbolic orbits, but these clearly originate in the solar system and only appear unbound at present due to minor gravitational and non-gravitational perturbations, so are not considered ICs. Indeed, identification of an object as a \emph{bona fide} IC in the usual orbit determination process would be straightforward, since ICs would have a highly hyperbolic orbit, i.e., eccentricities clearly greater than 1. Future advanced sky surveys, particularly the Large Synoptic Survey Telescope (LSST), will be many times more sensitive than past or present observations \citep{LSST}. It is therefore natural to consider whether LSST will detect ICs. The motivation for finding ICs is two-fold: discovering an IC would provide new observational opportunities and an IC would be an \emph{in situ} sample of another solar system. Like the discovery of the population of asteroids $\sim$200 years ago and the discovery of the Kuiper belt 20 years ago \citep{1993Natur.362..730J}, the eventual discovery of ICs will open new and unique avenues for exploration that will improve our understanding of the formation and evolution of planetary systems. Photometric, astrometric, and spectroscopic investigations can reveal estimates of the origin, physical, and chemical properties of a piece of another solar system. Even without detailed follow-up observations, the currently unknown frequency of ICs is a useful insight into the efficiency of planet formation in the galaxy. By estimating the frequency at which we expect to observe ICs and then comparing the expected value to the actual observational frequency, we can adjust planet formation models accordingly \citep[e.g.,][]{1990PASP..102..793S}. No matter the particulars, discovering ICs or placing upper limits on their frequency would help us to place our solar system in galactic context. However, as a rarely observed population with unique orbital properties, searching for ICs in LSST data will likely require a significant dedicated effort. The value of this effort depends partly on whether the frequency of ICs detected by LSST will have the power to discriminate between different planet formation models. For this reason, we provide a careful assessment of the sensitivity of LSST to different parameters of the IC population. The assessments of the frequency of detectable ICs has been highly variable ($\S$\ref{sec:back}). Earlier studies predicted very high numbers of observable ICs, usually by taking what was known from the formation of our solar system, estimating how many comets our solar system ejected into interstellar space, and then multiplying by the number density of stars. For example, \citet{McGlynn89} predicted that the number density of 1 km ICs was approximately $10^{13}$~pc$^{-3}$ and that this implied that several ICs should have already been detected. Recently, a careful assessment of the frequency of ICs by \citet[][hereafter M09]{M09} showed that the actual mass density is orders of magnitude less than these previous estimates, resulting in a number density for ICs larger than 1 km of $10^{5-10}$~pc$^{-3}$. M09 is the first study to self-consistently account for several realistic properties of the IC population by incorporating the stellar mass function, giant planet frequency estimates, solar system minor planet size distributions, and a more recent understanding of the formation of planetary systems. M09 considered the detectability of ICs by LSST, providing a clear explanation of why we have not observed any ICs to date. However, their analytical investigation was limited to considering ICs at the distance of Jupiter and beyond. They concluded that LSST would not be able to detect ICs at this distance. However, there are several aspects that may significantly enhance the frequency of detectable ICs closer than Jupiter: gravitational focusing by the Sun would enhance the concentration of ICs within Jupiter's orbit; ICs may brighten by several magnitudes due to outgassing; much more frequent smaller bodies can be seen at closer distances in a magnitude limited survey; etc. In order to estimate the realistic detectability of interstellar comets, we have developed a numerical simulation in order to consider all of the factors that play a role in detecting ICs. Our model includes effects from: \begin{itemize} \item increased density due to gravitational focusing (the effect of the Sun altering the trajectory of ICs\index{gravitational focusing}); \item photometric phase functions (the effect of observing ICs at different angles \index{photometric phase angle}); \item comet brightening (accounting for the increase in brightness of comets as they approach the Sun); \item conditions required for observability such as solar elongation (the angular distance between the IC and the Sun) and air mass (e.g., constraints on the altitude of topocentric observations). \end{itemize} These realistic factors will be discussed in full detail in $\S$\ref{sec:methods}. In addition, our method allows for a determination of many other IC properties relevant to observers, such as typical orbital parameters, rates of motion, and sky distribution. It is worth noting that the results of this paper can be divided into two parts. The orbit propagation and astrometry is based on some small assumptions, but is mostly robust. The estimation of the number of ICs that could be detected by LSST, on the other hand, requires several assumptions, in some cases using quantities not known even to within an order of magnitude which we leave as tunable parameters. In this regard, we occasionally neglect effects that would change the highly uncertain results by a factor of $\sim$2. The large uncertainty in our estimates should not be seen as a drawback of the model, but rather a motivation to search for ICs in order to place constraints on their currently unknown properties, with implications for planet formation theory.

\label{sec:discussion} The likelihood of detecting interstellar planetesimals has had a long and varied history. A modern understanding of the properties of the IC population was recently proposed by M09. In this work, we've studied the realistic observational aspects of detecting and characterizing this unknown population. Using a numerical model that tracks the position and brightness of ICs, we estimate that LSST could detect on the order of 1 IC during its 10 year lifetime, with orders of magnitude uncertainty mostly based on the actual frequency of small ICs. The expected size distribution of objects reduces this rate to $\sim$0.001, but including the contribution from interstellar asteroids or comet outbursts or discrete sources may boost the detection by 1-2 orders of magnitude. Frankly, some optimism is required to conclude that LSST will detect even 1 interstellar object. While it is possible to improve our model, our results are sufficiently informative to begin the discussion of whether and how the astronomical community should conduct the search for ICs. Facing the stark realization that ICs are exquisitely rare, we can expect to find them at the threshold of detection. In a single night, they are generally indistinguishable from NEOs, asteroids, and long-period comets. Only at the time of the next LSST observations (nominally 3 and 6 days later), will it become clear that the orbit can only be fit when the eccentricity is greater than 1. They are moving rapidly ($\sim$200 arcsec/hr, $\sim$1 deg/day) and will be difficult to link between single night detections. It is also likely that there will be occasional false positive linkages between unrelated solar system bodies that initially appear to be ICs. Algorithms attempting to detect solar system bodies may choose to discard detections and/or linkages that indicate an unbound orbit. This may be an easy way to help make the solar system moving object search more tractable, even though it would throw away any ICs that could nominally be detected. If possible, we recommend that systems searching for NEOs, asteroids, and/or comets \citep[such as the Pan-STARRS Moving Object Processing System;][]{den2012a} refrain from explicitly or implicitly biasing their systems against the algorithmic detectability of ICs, though these are vastly less frequent than solar system small bodies. The Pan-STARRS MOPS, in particular, is not explicitly biased against hyperbolic orbits. ICs convey rare and unique planet formation information; rare because ICs are so hard to observe and unique because their observations complement other methods used to study planet formation. The work needed to discover ICs is accompanied by a strong desire to follow them up with other observations, both for orbit recovery and for detailed characterization (e.g., with JWST). For this reason and based on our simulations, the ideal case is to discover and track ICs within 1-4 weeks, similar to NEOs. However, one of the strongest pieces of information gained from discovering an IC is their frequency and this could be determined in a specialized \emph{post facto} search, well after any detected ICs are recoverable or observable. Given the significant probability that LSST will not detect any ICs, such a project should be prepared to place an upper limit on the IC frequency based on a null detection. While it seems difficult to imagine now, we look forward to the day --- perhaps in the distant future --- that ICs are detected in such abundance that, like KBOs and exoplanets, the number of objects rapidly grows from zero to one to ten to a population so large it is hard to keep track of individual objects. It is exciting to consider what this future regime of IC studies could reveal about the formation and evolution of planetary systems in the Galaxy.

1607.08162

1607

1607.04385_arXiv.txt

Sulfur and zinc are chemically volatile elements, which play significant roles as depletion-free tracers in studying galactic chemical evolution. However, regarding red giants having evolved off the main sequence, reliable abundance determinations of S and Zn seem to be difficult despite that a few studies have been reported so far. Given this situation, we tried to establish the abundances of these elements for an extensive sample of 239 field GK giants ($-0.8 \ltsim$~[Fe/H]~$\ltsim +0.2$), by applying the spectrum-fitting technique to S~{\sc i} 8694--5, S~{\sc i} 6757, and Zn~{\sc i} 6362 lines and by taking into account the non-LTE effect. Besides, similar abundance analysis was done for 160 FGK dwarfs to be used for comparison. The non-LTE corrections for the S and Zn abundances derived from these lines turned out $\ltsim$~0.1(--0.2)~dex for most cases and not very significant. It revealed that the S~{\sc i} 6757 feature is more reliable as an abundance indicator than S~{\sc i} 8694--5 for the case of red giants, because the latter suffers blending of unidentified lines. The finally resulting [S/Fe]--[Fe/H] and [Zn/Fe]--[Fe/H] relations for GK giants were confirmed to be in good agreement with those for FGK dwarfs, indicating that S and Zn abundances of red giants are reliably determinable from the S~{\sc i} 6757 and Zn~{\sc 6362} lines. Accordingly, not only main-sequence stars but also evolved red giant stars are usable for tracing the chemical evolution history of S and Zn in the regime of disk metallicity by using these lines.

Sulfur and zinc are counted as important elements in stellar spectroscopy. Owing to their ``volatile'' characteristics, they are hardly affected by condensation process (unlike refractory elements such as Fe, which may suffer gas--dust depletion). For this reason, S ($\alpha$ group) and Zn (Fe-group element) are regarded to play significant roles in galactic chemical evolution, since they may be used as depletion-independent tracers of nucleosynthesis history. Accordingly, not a few spectroscopic investigations on the behavior of [S/Fe] and [Zn/Fe] in stars of various metallicities have been done toward clarifying the chemical enrichment history of these elements in the Galaxy (see, e.g., Takada-Hidai et al. 2005 for S, Caffau et al. 2005 for S, Saito et al. 2009 for Zn, Takeda et al. 2005c for S and Zn, and the references therein). However, regarding stars of Galactic disk population, unevolved dwarfs of F--G type have been primarily used for such studies of S and Zn abundances, while few reliable abundance results of these elements are available for red giants (evolved low--to-intermediate-mass stars), despite that several extensive abundance studies on a large sample of field giants have been reported in the past decade:\\ --- S and Zn were out of the target elements in Takeda, Sato, and Murata's (2008; hereinafter referred to as Paper~I) analysis for 322 G--K giants due to the limited wavelength range of the spectra.\\ --- In Luck and Heiter's (2007) spectroscopic study of 298 field giants, S and Zn abundances were determined only for 10 and 6 stars, respectively, which were not sufficient for establishing the trends of [S/Fe] or [Zn/Fe].\\ --- Luck (2015) recently reported the abundance results for a very extensive sample of 1133 field G--K giants. While S and Zn abundances were determined for most stars, the dispersions of [Element/Fe] ratios are unusually large ($\sim \pm 0.4$~dex) only for these two elements (cf. Figure~5 therein), which means that they are unreliable. Actually, as noted by Luck (2015) himself, [S/Fe] systematically grows from $\sim 0$ ($T_{\rm eff} \sim 5500$~K) to $\sim 2$ ($T_{\rm eff} \sim 4000$~K), manifestly indicating an involvement of appreciable blending (with unknown lines) becoming progressively important toward later spectral types (see Figure~6 therein). Regarding Zn, Luck (2015) similarly remarked that ``The zinc lines are badly blended, especially at temperatures below 4800K, resulting in abundances of poor quality.'' \\ --- Very recently, Maldonaldo and Villaver (2016) published the abundances of 16 elements (including S and Zn) for a large sample of evolved stars and compared those of unevolved main-sequence stars. As seen from their Fig. A.1, [S/Fe] results for evolved stars are systematically higher than those for main-sequence stars, and the dispersion of [Zn/Fe] for giants is appreciably larger than that for dwarfs, which implies that larger errors may be involved in their abundances of these elements for red giants.\\ --- In Jofr\'{e} et al.'s (2015) comparison of their abundance results for giant stars with literature studies (cf. Fig.~7 therein), Zn abundances were compared with those of Maldonaldo, Villaver, and Eiroa (2013); but the consistency was not necessarily satisfactory. It is thus evident that abundance studies of these two elements for evolved red giant stars are still at quite unsatisfactory level, which are yet to be conducted by a careful analysis of appropriate spectral lines. Takeda et al. (2015; hereinafter referred to as Paper~II) investigated the oxygen abundances (along with those of C and Na) for 239 late G and early K giants by using the $R$- and $I$-band spectra (up to $\sim 8800$~$\rm\AA$) newly obtained in the 2012--2013 period. This was originally intended to supplement/revise the previous work of Paper~I, which was based on old spectra covering only the yellow--orange region ($\sim$~5000--6200~$\rm\AA$). Since these new longer-wavelength region spectra comprises important spectral lines usable for determinations of S and Zn abundances (S~{\sc i} 8693--5, S~{\sc i} 6757, and Zn~{\sc i} 6362), we decided to make use these data and determine the S and Zn abundances of 239 red giants by using the spectrum-fitting method while properly including the non-LTE effect. Accordingly, we would like to examine in this investigation whether abundances of S and Zn can be reliably determined for low-gravity giants of G--K type. Toward this purpose, the S and Zn abundances of unevolved 160 dwarfs (including subgiants) of late-F through early-K type were also derived in essentially the same manner, which we intend to use as reference samples for comparing the abundance trends. That is, given that both S and Zn should not suffer any surface abundance change (unlike the abundances of some light elements such as Li or C, which are affected by evolution-induced mixing of nuclear-processed materials), we may expect that essentially the same [S/Fe]--[Fe/H] or [Zn/Fe]--[Fe/H] relation would be obtained for both GK giants and FGK dwarfs. If such a consistency could be confirmed, we may state that trustworthy abundances of these elements are established also for giants. Since abundance studies of two samples (GK giants and FGK dwarfs) are involved in this investigation based on different data sets by making use of our previous results, we firstly outline the adopted observational data, the atmospheric parameters, and the method of line-stregth measurement (to be described in section~2 and section~3) in table~1 for the reader's convenience. The remainder of this article is organized as follows: The abundance determination procedures (spectrum-fitting analysis, equivalent-width derivation, and derivation of non-LTE abundances) for our 239 program stars of GK giants are described in section~2. Section~3 is devoted to a brief description of S and Zn abundances for the reference sample of 160 FGK dwarfs. The resulting abundances along with the related quantities are discussed in section~4, where the [Fe/H]-dependent trends for giants and dwarfs are compared with each other. The conclusions are summarized in section~5.

Since sulfur and zinc belong to chemically volatile species, they play important roles in stellar spectroscopy because of being usable as depletion-free tracers for studying the galactic chemical evolution. However, regarding stars of disk population, spectroscopic studies of S and Zn abundances have been mostly directed to unevolved dwarfs. As a matter of fact, despite that several extensive abundance studies of evolved red giants have been published in the past decade, reliable abundance determinations for these two elements seem to be rare. Given this situation, we decided to establish the abundances of S and Zn for 239 apparently bright GK giants (in the metallicity range of $-0.8 \ltsim$~[Fe/H]~$\ltsim +0.2$) by applying the spectrum-fitting technique to S~{\sc i} 8694--5, S~{\sc i} 6757, and Zn~{\sc 6362} lines, where the non-LTE effect was also taken into account. Besides, S and Zn abundances for 160 FGK dwarfs were also determined in the same manner, which are to used for comparison. The non-LTE corrections for the S and Zn abundances derived from these lines turned out mostly $\ltsim$~0.1~dex (though amounting up to $\sim$~0.2--0.3~dex in some exceptional cases of S~{\sc i} 8695) and not very significant. The S~{\sc i} 6757 feature was found to be more reliable as an abundance indicator than S~{\sc i} 8694--5 for the case of red giants, because the latter appears to suffers appreciable blending with unidentified lines. The finally obtained [S/Fe]--[Fe/H] and [Zn/Fe]--[Fe/H] relations for GK giants were confirmed to be in good agreement with those for FGK dwarfs, indicating that S and Zn abundances of red giants are reliably determinable from the S~{\sc i} 6757 and Zn~{\sc 6362} lines. Consequently, not only unevolved main-sequence stars but also evolved red giant stars can be exploited as probe to study the chemical evolution history of S and Zn in the regime of disk metallicity. \bigskip Data reduction was in part carried out by using the common-use data analysis computer system at the Astronomy Data Center (ADC) of the National Astronomical Observatory of Japan. \newpage

1607.04385

1607

1607.03100_arXiv.txt

The Tully-Fisher relation (TFR) links the stellar mass of a disk galaxy, $M_{\rm str}$, to its rotation speed: it is well approximated by a power law, shows little scatter, and evolves weakly with redshift. The relation has been interpreted as reflecting the mass-velocity scaling ($M\propto V^3$) of dark matter halos, but this interpretation has been called into question by abundance-matching (AM) models, which predict the galaxy-halo mass relation to be non-monotonic and rapidy evolving. We study the TFR of luminous spirals and its relation to AM using the EAGLE set of $\Lambda$CDM cosmological simulations. Matching both relations requires disk sizes to satisfy constraints given by the concentration of halos and their response to galaxy assembly. EAGLE galaxies approximately match these constraints and show a tight mass-velocity scaling that compares favourably with the observed TFR. The TFR is degenerate to changes in galaxy formation efficiency and the mass-size relation; simulations that fail to match the galaxy stellar mass function may fit the observed TFR if galaxies follow a different mass-size relation. The small scatter in the simulated TFR results because, at fixed halo mass, galaxy mass and rotation speed correlate strongly, scattering galaxies along the main relation. EAGLE galaxies evolve with lookback time following approximately the prescriptions of AM models and the observed mass-size relation of bright spirals, leading to a weak TFR evolution consistent with observation out to $z=1$. $\Lambda$CDM models that match both the abundance and size of galaxies as a function of stellar mass have no difficulty reproducing the observed TFR and its evolution.

\label{SecIntro} The Tully-Fisher relation (TFR) links the luminosity of disk galaxies with their characteristic rotation speed. First noted by \citet{Tully-Fisher1977} using photographic magnitudes and HI velocity widths, it has become one of the best studied galaxy scaling relations and a powerful secondary distance indicator. It is well approximated, for luminous spirals, by a tight power law whose dependence on wavelength is fairly well understood \citep{Aaronson1979,Mathewson-Ford-Buchhorn1992,Verheijen1997,Tully1998,Haynes1999,Courteau2007}. As a result, the relation is now routinely cast in terms of galaxy stellar mass and the circular speed measured at a characteristic ``luminous radius'' \citep{Bell2001,Reyes2011,Avila-Reese2008}. Since rotation curves are nearly flat the choice of radius is not critical for luminous spirals, but popular choices include $2.2$ times the exponential scalelength \citep[e.g.,][]{Courteau1997} or, alternatively, a radius that contains roughly $80\%$ of all stars \citep[e.g.,][]{Pizagno2007}. The evolution of the TFR with redshift has been more difficult to pin down, although the consensus is that the TFR evolves weakly, if at all, up to $z \approx 1$. Early studies, many of them in the B-band, claimed significant evolution in either the zero-point, the slope or in both \citep[e.g.][]{Ziegler2002,Bohm2004}, but these conclusions evolved once data on longer wavelengths less affected by extinction became available. \citet{Conselice2005} and \citet{Flores2006}, for example, found no significant evolution in the K-band TFR to $z \approx 1.3$ and $z \approx 0.6$, respectively. This conclusion has been supported by the more recent work of \citet{Miller2011}, who conclude that there is no substantial change in the stellar-mass TFR out to redshifts of about unity. Observations at higher redshifts hint at more substantial evolution of the zero-point although the presence of large random motions and selection effects at such early times complicate the interpretation \citep{Forster-Schreiber2009,Cresci2009,Kassin2012}. The properties of the TFR have long challenged direct numerical simulations of disk galaxy formation in the $\Lambda$CDM scenario. Early work, for example, produced galaxies so massive and compact that their rotation curves were steeply declining and, at given galaxy mass, peaked at much higher velocities than observed \citep[see, e.g.,][and references therein]{Navarro2000,Abadi2003,Scannapieco2012}. The problem was quickly traced to the inability of early feedback schemes to prevent large amounts of low-angular momentum baryons from accumulating early at the center of dark matter halos. Subsequent work made progress by adopting feedback schemes able to remove a large fraction of the early-collapsing baryons and to regulate their further accretion, leading to disks with sizes and rotation curves in better accord with observation \citep[e.g.,][]{Okamoto2005,Governato2007,Brook2011,McCarthy2012,Guedes2013,Aumer2013,Marinacci2014}. Although such results were promising, they were inconclusive, especially because they were either based on a handful of carefully selected, and therefore likely highly biased, individual systems, or on cosmological boxes where simulated galaxies failed to match basic statistics of the observed galaxy population, such as the galaxy stellar mass function. As a result, much theoretical TFR work in the context of the $\Lambda$CDM cosmology has proceeded via semi-analytic models of galaxy formation. These models employ simple, albeit well-founded, prescriptions to generate a synthetic galaxy population within an evolving population of dark matter halos. The physical properties of such a population are then compared with observed galaxies in order to calibrate the assumed prescriptions and to shed light onto the role of various mechanisms during galaxy formation \citep[see, e.g.,][and references therein]{Cole2000,Dutton2010}. Semi-analytic models have highlighted a number of difficulties, particularly when attempting to match simultaneously the abundance of galaxies as a function of stellar mass and the slope and normalization of the TFR \citep[see][for recent attempts]{Lacey2015,Desmond2015}. The basic reason for these difficulties is that these models generally (and reasonably) assign more massive galaxies to more massive halos, leading to a tight relation between galaxy and halo masses that places strong constraints on their characteristic circular speed. A simple model for this galaxy-halo mass relation may be derived by ranking galaxies by mass and assigning them to halos ranked in similar fashion, preserving the ranked order \citep{Frenk1988,Vale-Ostriker2004,Guo2010,Moster2013,Behroozi2013}. This ``abundance-matching'' (AM) exercise has proven particularly useful when assessing the results of numerical simulations, especially those of single isolated systems, where there is otherwise little guidance about the mass or size of the galaxy that may form in one particular halo. Since the dark mass profile of $\Lambda$CDM halos is well known \citep{Navarro1996,Navarro1997}, AM models have little freedom left when trying to match the TFR: a galaxy's characteristic circular velocity is fixed once its radius and the halo response have been specified \citep[see, e.g.,][]{Cattaneo2014}. The critical role of galaxy size and halo response implies that insight into the origin of the TFR requires a good understanding of the interplay between galaxies' mass and size, as well as of the mass of the halos they inhabit and how galaxies might modify them. These complex issues are best studied through cosmological hydrodynamical simulations, especially those able to follow statistically significant numbers of galaxies over large volumes, and to resolve their inner regions, where rotation speeds are measured. These conditions are well met by the latest round of cosmological hydrodynamical simulations, such as the recently-completed Illustris and EAGLE projects \citep{Vogelsberger2013,Schaye2015}. One main conclusion from these efforts is that, except for the lowest masses \citep{Sawala2013,Sawala2015}, abundance-matching predictions are actually quite robust: matching the observed galaxy stellar mass function requires simulations to reproduce accurately the galaxy-halo mass relation envisioned by AM models, with little scatter. One intriguing result, however, is that both Illustris and EAGLE report good agreement with the observed TFR, despite the fact that the galaxy stellar mass functions they report differ significantly. This approximate ``invariance'' of the simulated TFR has been noted in the past. \citet{Guo2010}, for example, found that a number of simulated galaxies, which in earlier work had been reported to match the TFR, actually had masses that greatly exceeded AM predictions. A similar result has been discussed recently by \citet{Torrey2014}, who showed that the TFR in their simulations is insensitive to large variations in the Illustris galaxy formation physics submodules: only models with ``no feedback'' were found to be in substantial disagreement with the observed TFR. Although \citet{Torrey2014} cite ``feedback'' as an essential ingredient to match the TFR, its actual role in determining its slope and zero-point remains unclear, a point underlined by the recent results of \citet{Crain2015}, who report that the TFR is actually quite sensitive to feedback, at least in their implementation. We examine these issues here using the EAGLE set of $\Lambda$CDM cosmological hydrodynamical simulations. We analyze the stellar mass TFR in the regime of luminous spirals, where gas contributes little to the overall baryon budget, and report results on the baryonic Tully-Fisher relation of gas-dominated, fainter galaxies in a separate paper \citep{Sales2016}. We pay particular attention to the effect of galaxy sizes on the TFR, an issue that has been relatively well explored in semi-analytic approaches but that has received little attention in direct simulation TFR work. We begin in Sec.~\ref{SecModel} by motivating the effect of galaxy size on the TFR by simple considerations that highlight the need for halo contraction in order to reconcile the TFR with the predictions of abundance matching models. We then present, in Sec.~\ref{SecResults}, the TFR of simulated galaxies in EAGLE, with particular attention to the origin of its small scatter (Sec.~\ref{SecScatter}) and its evolution (Sec.~\ref{SecTFEVol}). We conclude with a brief summary of our main findings in Sec.~\ref{SecConc}. \begin{center} \begin{figure*} \advance\leftskip-0.85cm \includegraphics[width=200mm]{FigModel.png} \caption{Galaxy stellar mass, $M_{\rm str}$, as a function of various parameters. {\it Left:} The solid black curve shows the abundance-matching prediction of \citet[][B+13]{Behroozi2013}, as a function of halo virial velocity, $V_{200}$. Symbols correspond to the data of \citet[][P+07]{Pizagno2007}, converted to stellar masses using a constant I-band mass-to-light ratio of $1.2$ \citep[][]{Bell2003} and shown as a function of disk rotation speed, $V_{\rm rot}$. Color-shaded band indicates the mean slope and $1$-$\sigma$ scatter. {\it Middle:} Symbols show half-light radii of galaxies in the P+07 sample. Thick solid line indicates a multiple of $r_{\rm max}$, the characteristic radius where NFW halo circular velocities peak. Halo masses are as in the B+13 model of the left panel. {\it Right:} Tully-Fisher relation. The color band is the same as in the left-hand panel. The dotted curve indicates the dark halo circular velocity at $r_h=0.1\, r_{\rm max}$, assuming NFW profiles and neglecting the contribution of the disk. The dashed line includes the gravitational contribution of the disk, keeping the halo unchanged. Finally the thick solid line (and symbols) include the disk contribution {\it and} assume that halos contract adiabatically.} \label{FigModel} \end{figure*} \end{center} \begin{center} \begin{figure*} \advance\leftskip-0.1cm \includegraphics[width=180mm]{FigImages.pdf} \caption{Stellar surface density maps of three simulated disk galaxies at $z=0$. Stellar and halo masses, half-mass radii, and rotation parameter values are listed in the legends. The top row show a face-on view of the disks, the middle rows show edge-on views. The inner and outer circles indicate stellar half-mass radii, $r_h$, and $r_{\rm gal}=0.15\, r_{200}$, respectively. The corresponding circular velocity curves are shown in the bottom row. Blue denotes total circular velocity, grey the dark matter contribution, red the stars and orange the gas. Stellar half-mass radii and rotation speeds, $V_{\rm rot}=V_{\rm circ}(r_h)$, are indicated by dotted lines. The halo virial velocity, $V_{200}$, is shown with a horizontal arrow in each bottom panel. } \label{FigImages} \end{figure*} \end{center} \begin{center} \begin{figure*} \advance\leftskip-0.9cm \includegraphics[width=200mm]{FigSimTFR.png} \caption{Analogous to Fig.~\ref{FigModel}, but for EAGLE galaxies at $z=0$. Black solid lines are as in Fig.~\ref{FigModel}, and are included to aid comparison. Grey points correspond to all simulated galaxies, blue points indicate ``disks'' according to a relatively strict criteria; i.e., systems with rotation parameter $\kappa_{\rm rot}>0.6$. Galaxies forming in three narrow bins of halo mass are highlighted in cyan, red, and orange. The three starred symbols indicate the three galaxies shown in the images of Fig.~\ref{FigImages}. Note that EAGLE galaxies: (i) follow closely the B+13 abundance-matching predictions (left), (ii) have sizes comparable to spirals in the TF sample of P+07 (middle); and (iii) have a Tully-Fisher relation in good agreement with the predictions from the simple halo contraction model of Fig.~\ref{FigModel} (right). The thin grey line is a fit to the simulated TFR; see parameters in Table~\ref{TabFitParam}. } \label{FigSimTFR} \end{figure*} \end{center} \vspace*{-15mm}

\label{SecConc} We have used the EAGLE set of $\Lambda$CDM cosmological hydrodynamical simulations to study the relation between abundance matching, galaxy sizes, and the Tully-Fisher relation (TFR). Our main findings may be summarized as follows: \begin{itemize} \item Galaxies that match the predictions of abundance matching are consistent with the observed TFR despite the non-monotonic behaviour of the galaxy formation efficiency with halo mass. Consistency with the observed TFR requires galaxies to follow the mass-size relation of observed galaxy disks, and halos to respond to galaxy assembly by contracting, roughly as predicted by simple adiabatic contraction models. \item EAGLE galaxies match all of these constraints approximately, and show a Tully-Fisher relation in good agreement with observation at $z=0$. \item Galaxy size and halo contraction induce a strong correlation between galaxy formation efficiency and rotation speed that straightens the TFR into a single power law and scatters galaxies along the main relation, keeping its dispersion tight. As a result, the scatter of the simulated TFR is substantially lower than observed, even when {\it all} galaxies are considered, rather than only late-type disks. \item The EAGLE galaxy-halo mass relation evolves roughly as expected from AM models and its galaxy mass-size relation evolves in accord with that of galaxies in distant Tully-Fisher samples. This results in gradual but weak evolution of the simulated TFR in reasonable agreement with observed constraints, despite the more rapid evolution in galaxy formation efficiency predicted by abundance-matching models. \end{itemize} The sensitivity of the Tully-Fisher relation to galaxy size explains why many simulations have argued consistency with this scaling relation while, at the same time, failing to match the galaxy masses predicted by abundance-matching models. Indeed, it is always possible to trade disk mass for galaxy size so that the resulting TFR remains nearly invariant. A galaxy formation simulation cannot therefore be considered successful unless it matches simultaneously the Tully-Fisher relation, as well as the abundance and size of galaxies as a function of stellar mass. Overall, our results show that the slope, zero-point, scatter, and evolution of the TFR pose no obvious difficulty to galaxy formation models in the $\Lambda$CDM cosmogony.

1607.03100

1607

1607.01105_arXiv.txt

An approach for constructing a low-dissipation numerical method is described. The method is based on a combination of the operator-splitting method, Godunov method, and piecewise-parabolic method on the local stencil. Numerical method was tested on a standard suite of hydrodynamic test problems. In addition, the performance of the method is demonstrated on a global test problem showing the development of a spiral structure in a gravitationally unstable gaseous galactic disk.

During the last two decades, two main approaches have been used for numeral hydrodynamics simulations of astrophysical flows: the Lagrangian smooth particle hydrodynamics methods (hereafter, the SPH method) and Eulerian mesh-based methods. Numerous comparisons between the SPH and mesh-based methods have been performed \cite{Agertz_2007,Tasker_2008}. The main disadvantages of most SPH methods are the inaccurate computation of large gradients and discontinuities \cite{Vshivkov_2009}, suppression of physical instabilities \cite{Agertz_2007}, difficulty with choosing the proper smoothing kernel \cite{Attwood_2007} and the use of artificial viscosity \cite{Sijacki_2006}. The main disadvantages of most mesh-based methods are their Galilean non-invariance on the mesh \cite{Tasker_2008,Wadsley_2008}, difficulty with coding and implementation, and difficulty with treating multi-component systems, such as stars and gas \cite{Mitchell_2013}. During the last decade, the combined Eulerian-Lagrangian approach has been developed and actively used for numerical simulations of astrophysical hydrodynamics flows \cite{Murphy_2008,Springel_2010}. These methods unite the advantages of both approaches, while attempting to reduce the disadvantages. Recently, another Eulerian-Lagrangian numerical scheme for astrophysical flows has been presented, which employs the operator-splitting methodology and Godunov-type methods \cite{Vshivkov_2009,Kulikov_2013,Vshivkov_2007,Vshivkov_2011_a,Vshivkov_2011_b}. This numerical method is based on the solution of hydrodynamic equations in two stages. At the first Eulerian step, the hydrodynamic equations are solved excluding the advective terms. At the second Lagrangian step, equations for the advective transport are solved. This separation into two stages can effectively solve the problem of Galilean-invariance and the use of Godunov methods for the Eulerian step enables correct modelling of discontinuous solutions. This numerical method has been successfully employed for modeling collisions between galaxies \cite{Vshivkov_2011_a,Tutukov_2011}. {A significant disadvantage of the method lays} in the fact that it was first-order accurate. The main shortcoming of first-order schemes is dissipation of the numerical solution on discontinuities. It is therefore our main motivation to extend the original method to a higher order of accuracy. There are several well-known high-order numerical hydrodynamics methods such as the MUSCL (Monotonic Upstream-Centered Scheme for Conservation Laws) method \cite{KurganovTadmor_2000,VanLeer_1979}, total variation diminishing (TVD) method \cite{Jin_1995,Jiang_1996,Balsara_2000,Balsara_2009,Henrick_2005}, and piecewise parabolic method (PPM) \cite{Collela_1984}. The general idea of their approach is the construction of a piecewise-polynomial function on each cell of the numerical mesh. It may be a piecewise-linear reconstruction (in the case of the MUSCL scheme) or piecewise-parabolic reconstruction (in the case of the PPM method). For the construction of a monotonized numerical solution (needed to avoid the growth of spurious extrema) the limiters are usually employed in TVD methods \cite{Watersona_2007}. The problem of selection of limiters is analogous to the choice of artificial viscosity in SPH methods: the wrong choice of artificial viscosity for the SPH method (or limiters for the TVD method) can cause a substantial distortion of the numerical solution. The PPM method does not have the monotonicity problem, because piecewise-polynomial solutions on each cell are constructed without extrema. The main disadvantage of the PPM method is the use of a non-local stencil for computing hydrodynamic quantities on the next time step. The non-local stencil has problems with the proper choice of boundary conditions, domain decomposition, and dissipation of numerical solution. To resolve these problems, a modification of the PPM method was suggested -- the so-called piecewise-parabolic method on the local stencil (PPML) \cite{Ustyugov_2007,Ustyugov_2008}. The main idea of the PPML method is the use of a piecewise parabolic numerical solution on the previous time step for computing the Riemann problem. Numerical simulations of galaxies are an indispensable tool for studying their properties and evolution, taking into account the complexity of physical processes involved. The main galactic components such as gas and dust can be modelled as fluids, making numerical hydrodynamics simulations the main numerical approach. The main focus of this paper is concentrated on the extension of the original method \cite{Vshivkov_2009,Kulikov_2013,Vshivkov_2007,Vshivkov_2011_a,Vshivkov_2011_b} based on the PPML approach to a higher order of accuracy. {The methods presented here are specific to astrophysical flows where the extraordinarily high flow speeds, as well as the effect of self-gravity, impose further restrictions on the kinds of flow solvers that are designed. The operator splitting between the force terms and advected terms is certainly non-standard and would not be appropriate for a traditional flow solver.} In the first section, the details of the original method and the construction of a high order extension are described. The second section contains a verification of the high order nature of the modified method. In the third section, we provide computational experiments dealing with gravitational instability in disk galaxies and the development of a spiral structure.

{\bf The limiter problem.} The main trend in constructing accurate numerical methods (e.g. PPM, TVD and WENO) has been the use of high-order polynomials in the interpolation schemes. However, this practice often produces unphysical oscillations around discontinuous solutions. To resolve this problem, the limiters are introduced to create a monotonic interpolation scheme and remove nonphysical extrema \cite{Goloviznin_1998,Shu_1988,ShuOsher_1988}. However, the use of different limiters results often in the incorrect calculation of the shock wave velocity. The operator splitting approach allows us to use {limiters only when calculating piecewise parabolic functions in, e.g., equations~(33), (34), and (41)}. We do not provide a rigorous proof here, but we simply state that this property of our scheme has been experimentally confirmed. {\bf The Roe average.} In our numerical scheme, we provide a modification of the classic Roe space-averaging of hydrodynamic variables, which behaves better on modelling the gas-vacuum boundary. In fact, the use of the classic Roe scheme for density at the gas-vacuum boundary (and, in general, at any high density and pressure jumps) yields too large sound speeds. In general, the gas-vacuum boundary should be modelled by means a gas kinetic approach because of low collision frequency between gas particles, but that it is technically not feasible today. {\bf Advantages and disadvantages of our numerical method.} In our understanding, the advantages of our numerical scheme are: 1) accuracy on sufficiently smooth solutions and low dissipation on discontinuous solutions, 2) limiter-free and artificial-viscosity-free implementation, 3) Galilean invariance, 4) guaranteed non-decrease of the entropy \cite{GodunovKulikov_2014}, 5) extensibility on other hyperbolic models, 6) simplicity of program implementation, and 7) high scalability. We note that the use of a computational mesh results in distortion of the numerical solution. In the case of complex nonlinear hydrodynamic flows, such distortions can lead to the artificial alignment of the numerical solution with the cell boundaries. The use of non-regular or adaptive meshes alleviates but does not resolve this problem. Therefore, it is desirable to construct Galilean invariant numerical schemes, which produce numerical solutions that does not depend on the orientation of the computational mesh. We think that our numerical method can be effectively implemented on GPUs, as the GPUPEGAS code \cite{Kulikov_2014}, and on the Intel Phi architecture, as the AstroPhi code \cite{Kulikov_2015}, because it uses the same algorithms. Our numerical method has disadvantages, the three most significant of which we discuss below: \begin{enumerate} \item A rather simple finite-difference scheme for the time derivative. By example, the WENO code uses the Runge-Kutta fourth-order-accurate scheme for the time derivative. The advantage of such multilayer schemes is the expansion of the computational stencil (on regular 3D mesh to $9^3$ cells). It is a very strong solution for the Galilean invariance problem. \item The construction of the interpolation parabola for problems with complex geometries. When using non-regular meshes (e.g. triangle cells), the construction of piecewise parabola is a non-trivial problem, which requires the use of a special spline technique adapted to cells with an arbitrary simplex of cells. \item At the Eulerian step we solve the Riemann problem, which in turn relies on the analytic solution of the spectral problem (because the numerical solution of the eigenvalue/eigenvector problem is ill-posed). Finding the analytic solution for any hyperbolic equations is difficult. A possible solution is to use ''the potentials technique'' \cite{Godunov_2013}, but this technique requires the singular value decomposition of matrices, which is very expensive. \end{enumerate} {\bf The future work.} In the future, we plan to expand our numerical method to ''collisionless'' hydrodynamics, which solves for the first three moments of the collisionless Boltzmann equation \cite{Mitchell_2013,Vorobyov_2006}. The main features of this expansion will be the formulation of the equation of state for collisionless component and the development of thermodynamically consistent star formation process with guaranteed non-decrease of the entropy.

1607.01105

1607

1607.08738_arXiv.txt

The new CARMENES instrument is mounted at the 3.5 m telescope at Calar Alto Observatory, located in the Sierra de los Filabres in southern Spain. It consists of two fibre-fed high-resolution spectrographs, operating in the visible wavelength range from 0.52 to 0.96 $\mu$m and in the near-infrared from 0.96 to 1.71 $\mu$m, having a spectral resolution of R~>~80,000. \citep{Quirrenbach2010,Quirrenbach2012,Quirrenbach2014} Both spectrographs will simultaneously perform high-accuracy radial-velocity measurements of about 300 M dwarfs during three years of guaranteed observing time. The aim is to detect low-mass planets within the habitable zones of these stars. \\ For science preparation over 1500 high-resolution spectra have been observed with FEROS, CAFE and HRS to determine effective temperature, surface gravity and metallicity. These parameters are fundamental for characterising star-planet systems. The spectra of M dwarfs are very complex, with molecular lines forming due to the low temperatures. This makes it difficult to use a line-by-line approach and requires a full spectral synthesis, which in turn necessitates for accurate models that take into account the formation of molecules. We use the latest generation PHOENIX model grid, the PHOENIX ACES models \citep{Husser2013}. These models are especially designed for low temperature stellar atmospheres and use a new equation of state to accurately reproduce molecular lines.

We obtained stellar parameters for 351 stars from 977 spectra. We find that most stars lie within 3200-3900 K, corresponding to spectral types M1V-M5V, as shown in the upper left panel of Figure~\ref{fig:results}. The higher the metallicity the higher the temperature for each spectral type (Figure~\ref{fig:results}, lower left panel). This is consistent with results by \cite{Mann2015}. They showed that with increasing metallicity the radius increases, for fixed temperature. The spectral types have been calculated using spectral indices \cite{Schoefer2015}. The green squares correspond to a literature computation by \cite{PecautMamajek2013} for solar metallicity. A literature comparison with \cite{RojasAyala2012}, \cite{GaidosMann2014} and \cite{Maldonado2015} shows that our values for metallicity turn out to be higher than published ones. (Figure~\ref{fig:results}, upper right). One possible explanation for this is that PHOENIX ACES models still cannot reproduce the full depths of some lines (see Figure~\ref{fig:fit}, 4th wavelength range), which might cause the algorithm to choose higher metallicity models to fit the lines. On the other hand it seems that the signal-to-noise ratio is also very important for parameter determination. 75 percent of the stars with [Fe/H] higher than 0.6 have SNRs lower than 50. We find good agreement with expected [Fe/H] values for SNR>50 (Figure~\ref{fig:results}, lower right). For the first four months of CARMENES data we find that the parameters show better agreement with literature, having better SNRs. \begin{figure*} \centering \includegraphics[width=0.85\linewidth]{results.jpg} \caption{Temperature distribution of candidate sample (upper left), literature comparison for metallicity (upper right), spectral type-temperature relation (lower left, green dots: literature values for solar metallicity found by \cite{PecautMamajek2013}), metallicity distribution for stars observed with FEROS, CAFE and HRS for different SNRs (lower right).} \label{fig:results} \end{figure*}

1607.08738

1607

1607.05795_arXiv.txt

We explore the regulation of star formation in star-forming galaxies through a suite of high-resolution isolated galaxy simulations. We use the SPH code \textsc{Gasoline}, including photoelectric heating and metal cooling, which produces a multi-phase interstellar medium. We show that representative star formation and feedback sub-grid models naturally lead to a weak, sub-linear dependence between the amount of star formation and changes to star formation parameters. We incorporate these sub-grid models into an equilibrium pressure-driven regulation framework. We show that the sub-linear scaling arises as a consequence of the non-linear relationship between scale height and the effective pressure generated by stellar feedback. Thus, simulated star-formation regulation is sensitive to how well vertical structure in the ISM is resolved. Full galaxy disks experience density waves which drive locally time-dependent star formation. We develop a simple time-dependent, pressure-driven model that reproduces the response extremely well.

The process of star formation is limited, in principle, by the availability of cold, dense gas fuel. However, in typical disk galaxies it proceeds inefficiently relative to characteristic time-scales for the gas, such as the local free-fall time in Giant Molecular Clouds ($\sim 10$ Myr). In fact, globally, star formation proceeds on a time-scale comparable to several galactic rotation periods \citep[$\sim$1 Gyr,][]{kennicutt98,krumholzTan07}. This is commonly attributed to self-regulation. Highly resolved samples of nearby star-forming galaxies, such as The HI Nearby Galaxies Survey \citep[THINGS,][]{THINGSpaper}, the PdBI Arcsecond Whirlpool Survey \citep[PAWS,][]{PAWSpaper}, and the Panchromatic Hubble Andromeda Treasury \citep[PHAT,][]{PHATpaper} have provided a detailed look at the interstellar medium (ISM) on sub-kpc scales. This lets us connect the local physics of the ISM to the regulation of star formation on larger scales\citep[e.g.][]{krumholzMT09}. Although observations indicate that star formation must be regulated overall, the details of regulation are difficult to pin down using observational data due to the timing offset between tracers of gas and tracers of star formation \citep[e.g.][]{kruijssenLongmore14}. Thus the physical cycle of regulation is most easily studied theoretically. Simulations using isolated galaxies have tended to focus on exploring different kinds of feedback and overall star formation rates \citep{hopkinsQM11,hopkinsQM12}. Galaxy-scale simulations cannot directly resolve the detailed process of regulation, leading to increasingly complex sub-grid models. Studies of star-formation regulation have also been extended into the cosmological context \citep{agertz13,agertzKravtsov15}. Going to larger scales includes more of the galactic environment at the cost of lower resolution which potentially compromise elements of the regulation mechanism. \cite{ostriker10} presented a detailed semi-analytical model of equilibrium star formation and applied it to the THINGS galaxy sample. In this model, star formation is regulated by satisfying two equilibria in a galaxy. The first, a vertical dynamical equilibrium, requires a balance between ISM weight and pressure support. The second, a thermal or energy equilibrium, requires energy balance between feedback and heating/cooling processes \citep{ostriker10,ostrikerShetty11,kimKimOstriker11,kimOstriker15}. In this work we do not need to make any explicit assumptions regarding energy balance as the simulations do that explicitly. Following the first requirement, the ISM is compressed by its own weight, set by the gas column and the gravity of the various components of the galaxy (gas, dark matter and stars). The weight is then a function of the height of the gas layers. This sets an expected pressure at the mid-plane, which determines the cold gas fraction via the two-phase instability. Star formation occurs steadily in cold, dense gas. To keep this model consistent, the ISM must have sufficient means of support to provide that mid-plane pressure. In the \cite{ostriker10} model, this pressure is effectively set by the star formation rate. The equilibrium semi-analytic model provided an excellent fit to star formation rates in the THINGS galaxies, where the vertical gravity is dominated by stars. Determining the total effective pressure is key for vertical pressure balance. The dominant mechanisms providing this effective pressure support can change in different environments. In most nearby galaxies ($\Sigma_g \gtrsim 10$ M$_{\odot}$ pc$^{-2}$), however, pressure from turbulent support plays the major role. This model assumption has been tested on small scales and in two- and three-dimensional simulations \citep{kimKimOstriker11,ostrikerShetty11}. Stellar feedback is responsible for generating this support. This may come from supernovae, however, in starbursts, some authors argue that alternate mechanisms, such as radiation pressure, play a major role \citep{hopkinsQM11}. Supernovae as regulators of star formation have been explored through small-box simulations in several works \citep[e.g.][]{deAvillez05,joung09,creasey13}. In the prior scenarios, the effective pressure is strongly linked to the local star formation rate. We note in passing that in the outer regions of disks, ($\Sigma_{\text{g}} \lesssim 10$ M$_{\odot}$ pc$^{-2}$), gas could be supported due to non-local heating or turbulence associated with galactic shear \citep{mcnally09}. A two-phase structure may not develop, limiting star formation in these environments \citep{elmegreenParravano94, schaye04}. Local simulations with a fixed surface density are conducive to reaching a steady equilibrium and thus ideal to study the hydrostatic, equilibrium framework of the \cite{ostriker10} model. Pressure-driven regulation should also play a key role in disk-scale simulations, with the addition of time dependent, spatial variations. Isolated galaxy simulations naturally include all of the necessary components (e.g. large-scale shear) but still allow relatively high resolution for the ISM. High-resolution simulations of isolated galaxies have been extensively used to study the formation of dense star-forming gas \citep{dobbs08,dobbs11,renaud13,taskerTan09,tasker11}. However, a realistic feedback and star-formation regulation cycle was not a major focus of these works. Examining pressure balance and star-formation regulation requires high numerical resolution. Fundamentally, one must be able to resolve the scale height of the gas disk, particularly for the colder phases, where stars form, which have scale heights of order 100 pc or less. Forming multiple phases is also numerically demanding. Turbulence is an important contributor to the detailed structure of star forming clouds. As a result, galaxy-scale simulations can directly examine star cluster-scale formation at best. Thus even high-resolution simulations must still rely on star formation recipes. Turbulence also contributes to regulation and pressure support. Galactic turbulence is generated at a range of scales up to several kpc and then cascades down to smaller scales. Simulations have difficulty maintaining this cascade. In particular turbulent energy is typically suppressed (via numerical dissipation) on scales substantially above the resolution limit \citep{priceFederrath10}. In galaxy-scale simulations this severely limits turbulence on small scales in the ISM. However, turbulence on scales comparable to the scale height is strongly linked to stellar feedback and thus star-formation regulation. An attractive option is to avoid the numerical issues by injecting energy on small scales as an effective pressure as in \cite{agertz13}. The pressure-driven regulation framework has profound implications for simulators. The fact that most simulated models are able to regulate star formation suggests that just providing a source of effective pressure linked to young stars is sufficient for basic regulation. In addition, star formation and feedback models typically have a lot of parameter freedom to tune star formation rates to match expectations. Thus achieving regulation or even tuning it to match a narrow set of observations is not that remarkable. The pressure-driven picture should explain these results if we incorporate typical sub-grid models into the framework. For example, a common outcome from the prior work discussed above has been that simulated star formation rates do not scale linearly with variations in star formation and feedback parameters \citep[e.g.][]{hopkinsQM11}. Linear scaling in the star formation rate constant, for example, would be the naive expectation. This assumes that star formation and the ISM are weakly coupled. However, variations in key ISM properties such as scale height are likely to be artificially limited at low resolution. By moving to higher resolution than has typically been employed for many sub-grid models in use today, one can probe how robust these results are. The above issues with sub-grid star formation in whole-galaxy simulations undermine the common approach of using regulated star formation as the primary benchmark for a successful model. The key question raised is whether the regulation achieved in simulated galaxies is actually similar to that occurring in nature. An associated question is how much predictive power these models have when used outside the cases on which they were calibrated. The pressure regulation idea is extremely general and should apply to all steadily evolving galaxies, even as conditions and modes of feedback are varied. Small scale simulations have confirmed the utility of these ideas for local, relatively uniform conditions \citep[e.g.][]{kimKimOstriker11, kimOstrikerKim13}. We would like to apply this insight to understand simulations on the scale of a whole galaxy. To do so we should extend the framework to include the time dependent behaviour, (e.g. in response to density waves), expected in the broader galactic context. To do so, one needs to move from vertical hydrostatic equilibrium to a non-equilibrium pressure-driven framework. In the current work, we employ a suite of high resolution simulations of isolated galaxies to explore basic ideas of star-formation regulation. Though we use high resolution and incorporate a fairly complex ISM model, we have kept the star formation and feedback prescriptions simple to make the results easier to interpret. None-the-less, these models are quite similar to the most popular star formation and feedback models in use in simulations at the current time and should therefore also provide direct insight into how those models operate. Given that our resolution is much higher than much of the early work in which the sub-grid models were developed \citep[e.g.][]{stinson06}, a first step is to test basic sensitivity to parameter choices. We can also compare to other recent, high-resolution simulations of isolated galaxies, which have tended to use much more complex feedback \citep[e.g.][]{hopkinsQM11}. A further feature of this work is to test the pressure-driven framework in a dynamic environment with time dependence and spatial variations (e.g. density waves) arising naturally in a galactic setting. This necessitates extending the model into a time-dependent, (non-equilibrium) pressure-driven framework. The structure of the paper is as follows. In section~\ref{sec:method}, we describe our simulation set-up for an isolated galaxy and our simple feedback recipe. In section~\ref{sec:results}, we characterize the overall simulated galaxy behaviour. In particular, it shows sub-linear scaling of star formation rates with respect to the model parameters and we explore how these arise. In section~\ref{sec:balance}, we demonstrate that the galaxies were behaving in a manner consistent with expectations from the pressure balance framework. Section~\ref{sec:timevary} demonstrates the local time variability in star formation and associated ISM properties. Finally, in section~\ref{sec:dynamicpressure}, we present a time-dependent model extension of the pressure-driven framework. This reproduces the behaviour seen in the simulations including the sub-linear scaling with star formation parameters. \begin{figure} \centering \includegraphics[scale=0.42]{figure1.pdf} \caption{Phase diagram for the reference galaxy at a time of 500 Myr. Gas was sampled in annuli centred at the labelled radii +/- 50 pc, with different colouring denoting different radii. Here, darker colouring denotes more mass residing at that location in the phase space. Since the photoelectric heating term varies with radius, the equilibrium is different at different radii in the disk. This variation changes the minimum density required to maintain a two-phase structure in the ISM.} \label{fig:phase} \end{figure}

\label{sec:conclusions} We have presented a suite of high resolution isolated galaxies that include a purposefully simple feedback scheme. We use a wide range of parameter choices to explore both the implications of these choices and the process of regulation of star formation. We find a sub-linear scaling between parameter choices and the resulting amount of star formation. This is found in many other simulations \citep[see, e.g.][]{hopkinsQM11}. We adapted the equilibrium pressure regulation models of \cite{ostriker10} and \cite{kimKimOstriker11, kimOstrikerKim13} to include representative simulation sub-grid recipes for star formation and feedback. This allowed us to demonstrate the origin of the sub-linear scaling with these parameters. The equilibrium models readily explain how feedback affects the star formation rates. The effect of changing the overall star formation efficiency is also well represented by the models, however, strongly non-linear parameters, such as star formation density thresholds, are harder to incorporate. Such parameters also reflect the complexity of star-formation and how difficult it is to model even in three-dimensional simulations. The simulations show regular variations in the local star formation rates and other properties. These are driven by density waves which occur naturally in global models. In comparison, small box simulations \citep[e.g.][]{kimOstrikerKim13} are expected to approach rough equilibrium and do not need to consider vertical oscillations. We extended the equilibrium models into a dynamic pressure-driven regulation framework. We show that advection plays a minor role and thus the vertical motions are driven by local effective pressure differences. We adapt the star formation and feedback models used in the simulations to produce a complete local model of star formation and feedback. These models are able to qualitatively reproduce the behaviour of our simulated galaxies, including the variability and the sub-linear scaling with the star formation and feedback efficiency parameters. A goal of this work was to determine which aspects of the small-scale star formation physics strongly affect larger scale simulations. The simulation community has invested in a diverse array of feedback prescriptions and types. We demonstrate that the crucial factor is how feedback translates into effective pressure support on larger scales. It is the effective pressure support that regulates star formation and the vertical structure of the ISM. Realistic regulated star formation thus requires that the scale height be resolved. Variable FUV backgrounds are a potentially important stellar feedback for normal spiral galaxies. Including this variability is feasible in local boxes \citep{kimOstrikerKim13} but is numerically challenging in global galaxy simulations. In future work, we plan to explore this mode of feedback using newly developed radiative transfer techniques (Woods et al., in preparation) including an old stellar disk. The inclusion of FUV that is coupled to star formation allows self-consistent measurement of the partition between different pressure components \citep{kimOstriker15,koda16}.

1607.05795

1607

1607.00386_arXiv.txt

We build a new model for the global 21-cm signal that is calibrated to measurements of the high-$z$ galaxy luminosity function (LF) and further tuned to match the Thomson scattering optical depth of the cosmic microwave background, $\tau_e$. Assuming that the $z \lesssim 8$ galaxy population can be smoothly extrapolated to higher redshifts, the recent decline in best-fit values of $\tau_e$ and the inefficient heating induced by X-ray binaries (the presumptive sources of the high-$z$ X-ray background) imply that the entirety of cosmic reionization and reheating occurs at $z \lesssim 12$. In contrast to past global 21-cm models, whose $z \sim 20$ ($\nu \sim 70$ MHz) absorption features and strong $\sim 25$ mK emission features were driven largely by the assumption of efficient early star-formation and X-ray heating, our new models peak in absorption at $\nu \sim 110$ MHz at depths $\sim -160$ mK and have negligible emission components. Current uncertainties in the faint-end of the LF, binary populations in star-forming galaxies, and UV and X-ray escape fractions introduce $\sim 20$ MHz ($\sim 50$ mK) deviations in the trough's frequency (amplitude), while emission signals remain weak ($\lesssim 10$ mK) and are confined to $\nu \gtrsim 140$ MHz. These predictions, which are intentionally conservative, suggest that the detection of a 21-cm absorption minimum at frequencies below $\sim 90$ MHz and/or emission signals stronger than $\sim 10$ mK at $\nu \lesssim 140$ MHz would provide strong evidence for ``new'' sources at high redshifts, such as Population III stars and their remnants.

\label{sec:Introduction} Galaxy evolution in the early Universe is most often studied from two distinct vantage points: through direct observations of high-$z$ galaxies and measurements of the thermal and ionization history of the intergalactic medium (IGM). These two approaches are exceptionally complementary, at least in principle, as the IGM is a repository of photons that never reach our telescopes. At the highest redshifts, due to limitations of even the most powerful space-based optical and near-infrared instrumentation, future constraints on the properties of the IGM will serve as an essential substitute for direct observations of galaxies themselves. As a result, the establishment of a framework for inferring galaxy properties from IGM signals is paramount. The canonical probe of high-$z$ galaxies is the rest-frame ultraviolet (UV) galaxy luminosity function (LF), i.e., the number density of galaxies per unit luminosity and redshift. Dedicated programmes using the \textit{Hubble Space Telescope} (\textit{HST}) have driven progress in this area at the highest redshifts so far probed, with healthy samples now extending to redshifts as high as $z \sim 8$ \citep{Bouwens2015,Finkelstein2015}, and a number of candidates at $9 \lesssim z \lesssim 12$ \citep[e.g.,][]{Ellis2013, Oesch2014}. The \textit{James Webb Space Telescope} (\textit{JWST}) will be very important for filling out the sample of galaxies at yet higher redshifts, but current models predict its reach will not extend beyond $z \sim 15$ \citep[e.g.,][]{Mason2015} without the aid of strong lensing, which has boosted \textit{HST}'s capabilities in the Frontier Fields \citep{Atek2015,Livermore2016}. Complementary IGM-based constraints on high-$z$ galaxies are far more crude at this stage. The Thomson Scattering optical depth, $\tau_e$, for example, constrains the total column density of electrons between the observer and the cosmic microwave background (CMB), while Gunn-Peterson troughs in quasar spectra mark the end of reionization at $z \sim 6$ \citep[e.g.,][]{Fan2002}. Together, these constraints provide a lower limit on the duration of the Epoch of Reionization (EoR), which one can parameterize as a reionization redshift, $z_{\mathrm{rei}}$, assuming an instantaneous transition from neutral to ionized. The \textit{Planck} team recently reported a Thomson scattering optical depth to the cosmic microwave background (CMB) of $\tau_e = 0.055 \pm 0.009$ \citep{Planck2016}, indicating $z_{\mathrm{rei}} \sim 8 \pm 1$, and thus a minimal duration of $\Delta \zrei \sim 2$. Efforts to jointly interpret the aforementioned measurements have largely been geared toward reconciliation. Do the number of photons generated by the galaxies we do see match the number required to maintain a state of full ionization in the IGM? Furthermore, is the population of galaxies in place prior to full reionization substantial enough to match the most up-to-date measurements of $\tau_e$? The answer to both of these questions is ``yes,'' provided one makes reasonable assumptions about (i) the abundance of galaxies we do \textit{not} see (i.e., extrapolations to the LF), and (ii) the escape fraction of Lyman continuum (LyC) photons from galaxies, $\fescLyC$. Recent work suggests that extrapolating the Schechter form of the LF to low luminosities may well be reasonable \citep{Livermore2016}, at least at $z \sim 6$, while the $\fescLyC \sim 0.2$ values that have caused discomfort in recent years may be now reasonable if the UV emission of star clusters is boosted by binary star evolution \citep{Eldridge2009,Stanway2016,Ma2016}. Despite such reduced tensions between theory and observation, the story of galaxy evolution in the early Universe is far from complete. In the coming years, observations of redshifted 21-cm emission from neutral hydrogen are expected to contribute substantially to our existing understanding of reionization and high-$z$ galaxies while opening up a brand new window into the excitation (or ``spin'') temperature history of the high-$z$ IGM \citep[e.g.,][]{Madau1997,FurlanettoOhBriggs2006}. As a result, 21-cm measurements promise to weigh in on long-standing questions regarding the ionizing photon production efficiency of high-$z$ galaxies, the nature of X-ray sources \citep[e.g.,]{Pritchard2007,Fialkov2014,Pacucci2014,EwallWice2016}, and perhaps even the properties of the interstellar medium of high-$z$ galaxies, which can serve as both a source of radiation \citep[e.g., bremsstrahlung;][]{Mineo2012b} and sink for LyC, X-ray, and perhaps even Lyman-Werner photons \citep{Kitayama2004,Schauer2015}, whose escape fraction we consider in Section \ref{sec:escape} for the first time in a 21-cm context. Studies aimed at better-understanding the complementarity of 21-cm measurements and other EoR probes, though few so far, demonstrate that even crude 21-cm constraints can greatly aid our understanding of reionization and high-$z$ galaxies \citep{Pritchard2010b, Beardsley2015}. Preliminary results from the \textit{Precision Array for Probing the Epoch of Reionization} (\textit{PAPER}) have since bolstered these arguments, finding the first observational evidence of X-ray heating of the high-$z$ IGM through upper limits on the 21-cm power spectrum \citep{Parsons2014,Ali2015,Pober2015,Greig2016}, and thus constrained the X-ray properties of $z \sim 8$ galaxies for the first time. The complementarity can be viewed from the opposite perspective as well, since constraints on high-$z$ galaxies can in principle be used to better separate signal from foreground \citep{Petrovic2011}. The sky-averaged (``global'') 21-cm signal \citep{Shaver1999}, now being targeted by several ground-based experiments \citep[e.g., EDGES, BIGHORNS, SCI-HI, SARAS, LEDA;][]{Bowman2010,Sokolowski2015,Voytek2014,Patra2015,Bernardi2016}, with more concepts in design \citep[e.g., DARE;][]{Burns2012}, is a particularly clear-cut ally of galaxy surveys as it is sensitive to the volume-averaged (i.e., luminosity function integrated) emissivity of galaxies. The mean ionization and spin temperature histories encoded by the global 21-cm signal of course influence the 21-cm power spectrum as well. The joint constraining power of the power spectrum and global signal was recently considered by \citet{Liu2016b}, though to the best of our knowledge the 21-cm signal (sky average or power spectrum) and galaxy LF have yet to be considered in a common framework. This has prevented 21-cm models from calibrating to recent advances driven by \textit{HST}, and as a result has led to predictions spanning the a wide range of possibilities \citep[e.g.,][]{Furlanetto2006,Pritchard2010a,Mesinger2013,Fialkov2014,Mirocha2015,Tanaka2016}. Our goal here is to address these issues in two steps: \begin{enumerate} \item Leverage the success of simple models for the galaxy LF \citep[e.g.,][]{Trenti2010,Tacchella2013,Sun2016} to create a new ``vanilla model'' for the global 21-cm signal calibrated both to the LF \citep{Bouwens2015} and $\tau_e$ \citep{Planck2016}. \item Explore simple extensions to the standard picture of the LF in an attempt to determine the global 21-cm signal's sensitivity to the properties of the faint galaxy population, and thus more concretely determine how its detection will complement future galaxy surveys and 21-cm power spectrum experiments. \end{enumerate} In the near term, these models can be used to test signal extraction algorithms and better inform instrument design. In the longer term, our models will provide a reference point from which to interpret a global 21-cm measurement in the broader context of galaxy formation. This paper is organized as follows. In Section \ref{sec:model}, we outline our theoretical model for the galaxy population and global 21-cm signal. In Section \ref{sec:results}, we present our main results, including our LF-calibrated model for the global 21-cm signal, its sensitivity to the star-formation efficiency of faint galaxies, the stellar and black hole populations of high-$z$ galaxies, and the escape fraction of UV and X-ray photons. We provide some discussion of our results in Section \ref{sec:discussion} and summarize our main conclusions in Section \ref{sec:conclusions}. We use cosmological parameters from \citet{Planck2015} throughout.

\label{sec:conclusions} We have shown that linking models of the global 21-cm signal to recent measurements of the high-$z$ galaxy LF leads to a strong preference for models with late heating and reionization ($z \lesssim 12$). At the formalism level, we assume that the $\fstar$-based model is the true model for the galaxy luminosity function, that the stellar population synthesis models we employ are accurate, and that the IGM is reasonably well-modeled as a two-phase medium. As for the components of the model, we assume that high-mass X-ray binaries are the sole sources of the $z \gtrsim 6$ X-ray background, and follow a $L_X$-SFR relation similar to that of local star-forming galaxies. If these assumptions hold, then our model suggests that: \begin{enumerate} \item The global 21-cm signal peaks in absorption at $\nu \sim 110$ MHz and a depth of $\sim -160$ mK. The emission feature is negligible in most models, reaching an amplitude $\lesssim 10$ mK at $\nu \sim 150$ MHz in the most optimistic scenario (Figures \ref{fig:gs_fiducial} and \ref{fig:gs_Z}). Ruling out such models may be easier than the sharp step-function emission models typically targeted at $\nu \gtrsim 100$ MHz, and would provide clear evidence of non-standard physics and/or source populations. \item The absorption trough in the global 21-cm signal is very sensitive to the SFE in low-mass galaxies (Figure \ref{fig:gs_fiducial}), with a $\sim 30$ MHz spread in its position arising from differences between currently viable models. Constraining the SFE will be very difficult with the LF alone given the depths one must probe ($\MUV \sim -12$) to differentiate models. \item The minimum mass of halos capable of supporting star formation, parameterized through $\Tmin$, has only a minor impact on our results (Figure \ref{fig:Tmin_effects}) given the steep decline of the SFE with halo mass implied by the LF (Figure \ref{fig:calib_sfe}). \item The $Z$-dependence of the $L_X$-SFR relation is very important, affecting the depth of the absorption feature at the $\sim 50$ mK level (Figure \ref{fig:gs_Z}), which corresponds to mean IGM spin temperatures between $\sim 10$ and $40$ K at $z=8.4$ (Figure \ref{fig:Ts_QHII_tau_Z}), close to the recent \textit{PAPER} constraints. Stellar metallicity plays a relatively minor role in setting the ionization history (Figure \ref{fig:Ts_QHII_tau_Z}). \item The global 21-cm signal is not very sensitive to the LyC emission properties of galaxies. However, the escape of LW and X-ray photons may have a dramatic impact on the signal (Figure \ref{fig:escape}). \end{enumerate} J.M. would like to thank Louis Abramson for many stimulating conversations and comments on an earlier draft, and the anonymous referee for comments that helped improve this paper. This work was supported by the National Science Foundation through award AST-1440343 and by NASA through award NNX15AK80G. We also acknowledge a NASA contract supporting the ``WFIRST Extragalactic Potential Observations (EXPO) Science Investigation Team'' (15-WFIRST15-0004), administered by GSFC. SRF was partially supported by a Simons Fellowship in Theoretical Physics and thanks the Observatories of the Carnegie Institute of Washington for hospitality while much of this work was completed.

1607.00386

1607

1607.07874_arXiv.txt

We study a resistive tearing instability developing in a system evolving through the combined effect of Hall drift in the Electron-MHD limit and Ohmic dissipation. We explore first the exponential growth of the instability in the linear case and we find the fastest growing mode, the corresponding eigenvalues and dispersion relation. The instability growth rate scales as $\gamma \propto B^{2/3} \sigma^{-1/3}$ where $B$ is the magnetic field and $\sigma$ the electrical conductivity. We confirm the development of the tearing resistive instability in the fully non-linear case, in a plane parallel configuration where the magnetic field polarity reverses, through simulations of systems initiating in Hall equilibrium with some superimposed perturbation. Following a transient phase, during which there is some minor rearrangement of the magnetic field, the perturbation grows exponentially. Once the instability is fully developed the magnetic field forms the characteristic islands and X-type reconnection points, where Ohmic decay is enhanced. We discuss the implications of this instability for the local magnetic field evolution in neutron stars' crusts, proposing that it can contribute to heating near the surface of the star, as suggested by models of magnetar post-burst cooling. In particular, we find that a current sheet a few meters thick, covering as little as $1\%$ of the total surface can provide $10^{42}~$erg in thermal energy within a few days. We briefly discuss applications of this instability in other systems where the Hall effect operates such as protoplanetary discs and space plasmas.

A plethora of observations of strongly magnetised neutron stars \citep{Olausen:2014} has revealed that their temperatures are higher than what conventional cooling of a hot proto-neutron star suggests. A solution to this puzzle is that the extra thermal energy needed for these systems is provided by the Ohmic decay of their magnetic energy reservoir \citep{Pons:2007}. However, given the high conductivity of a neutron star crust, the rate of Ohmic decay is expected to be slow and the conversion of magnetic energy to heat inefficient. This has led to the idea that the Hall effect may be able to accelerate magnetic field decay, as the Hall timescale is inversely proportional to the intensity of the magnetic field. This acceleration can only be done in an indirect way, as the Hall effect conserves magnetic field energy. Several paths have been proposed in this direction. \cite{Goldreich:1992} suggested that the Hall effect may lead to the formation of smaller scale structure through cascades, which have reduced Ohmic decay times, a result that has been followed up by numerical studies exploring Hall-induced turbulence \citep{Biskamp:1996, Wareing:2009b, Wareing:2010}. Another possibility is the development of instability of a state previously being in Hall equilibrium leading to smaller structure formation \citep{Rheinhardt:2002, Rheinhardt:2004, Pons:2010}. Recent work of \cite{Wood:2014} found a family of exact solutions for the density-shear instability in electron-MHD, requiring a covarying magnetic field and electron number density, a result that was studied numerically in detail by \cite{Gourgouliatos:2015b}. Apart from instabilities and cascades, secular Hall evolution has been explored: \cite{Vainshtein:2000} studied the effect of the sharp drop of electron number in the crust, finding that the magnetic field evolution is described by a Burger's type equation, leading to the formation of shocks in the form of current sheets decaying on a Hall timescale rather than the slower Ohmic, and applied to the evolution of a toroidal field in an axially symmetric system by \cite{Reisenegger:2007}. Once the poloidal field is included \citep{Hollerbach:2002, Hollerbach:2004} the formation of current sheets is followed by an oscillatory behaviour. The consensus of axially symmetric crustal simulations, exploring a broad range of initial conditions \citep{Pons:2009, Kojima:2012,Vigano:2012,Gourgouliatos:2014a}, has concluded that the Hall effect drastically changes the structure of the magnetic field, whereas later, Hall evolution saturates \citep{Gourgouliatos:2014b}. An intrinsic drawback of global neutron star simulations is the fact that they under-resolve current sheets. Current sheets form both in the uniform electron density case \citep{Wareing:2010} and even more efficiently in the presence of an electron density gradient \citep{Vainshtein:2000, Vigano:2012}. Furthermore, they are likely to appear near the surface of the crust, as the available electric charges decrease dramatically from the solid crust to the plasma magnetosphere. In the latter case, a usual assumption made in simulations is that the external magnetic field is a vacuum potential field which leads to boundary effects by matching the two configurations \citep{Wood:2014}. In their seminal paper \cite{Furth:1963} studied finite-resistivity instabilities of a sheet pinch finding the so-called tearing instability {\it ``a long-wave `tearing' mode, corresponding to a breakup of the layer along the current flow lines"}. Linear analysis of the MHD system yields an exponential growth rate $\gamma_{T} \sim \tau_{O}^{-3/5}\tau_{A}^{-2/5}$, where $\tau_{O}$ and $\tau_{A}$ are the resistive and Alfv\`en times respectively, while in the non-linear phase the growth becomes algebraic \citep{Rutherford:1973}. Several applications of the tearing instability have been considered in astrophysical contexts. \cite{Rosenbluth:1967} studied resistive instabilities in magnetospheric tails. \cite{Priest:1985} presented various applications of the tearing instability in relation to current sheets developing in solar and space plasmas. The tearing instability is considered to be an efficient mechanism for powering solar flares and accelerating particles therein \citep{Sturrock:1966, Somov:1989, Aschwanden:2002}. Recent numerical simulations by \cite{Landi:2015, DelZanna:2016} in general astrophysical contexts have demonstrated the development of the tearing instability in the limit of very high conductivity for appropriately thin current layers. Other applications have focused on pulsar magnetospheres, where numerical simulations agree on the presence of current sheets, either confined to the equatorial plane as is the case in axially symmetric systems \citep{Contopoulos:1999, Komissarov:2006}, or with more complicated geometries for the case of inclined systems \citep{Spitkovsky:2006, Kalapotharakos:2009}. In depth study of the current sheets of pulsar magnetospheres by \cite{Uzdensky:2014}, showed that they are susceptible to the tearing mode instability leading to the formation of plasmoids with the eventual emission of high energy radiation and non-thermal particles \citep{Sironi:2014}. The tearing instability has also been studied in the context of Relativistic MHD considering applications to magnetar flares and jets through explosive reconnection \citep{ Komissarov:2007, Elenbaas:2016, Barkov:2016}. Motivated by the omnipresence of the tearing instability in current sheets and their formation in neutron star crusts through the Hall effect, we study its development and impact. We explore the evolution of the magnetic field in a configuration where the tangential component changes direction by $180^{\circ}$ within a thin layer, allowing for some finite resistivity, in the inertialess electron-MHD formulation. We show, through linear and non-linear calculations, that the tearing mode instability naturally appears and enhances the decay of the magnetic field. We note that the term Hall evolution (or drift) has the meaning of Electron-MHD when used to describe the evolution of the magnetic field in the crust of neutron stars. There, only electrons are allowed to moved through a solid crystal lattice consisting of positively charged ions \citep{Jones:1988}. In principle, Hall evolution can accommodate for the motion of more than one charged species whereas Electron-MHD refers to systems where only electrons move, making the latter a special case of the former. In this paper the term Hall-MHD is used in the limit of Electron-MHD. The plan of the paper is as follows: In Section 2 we formulate the equations of Electron-MHD. We solve these equations in the linear and non-linear regime in Section 3. We discuss the properties of the instability and compare it with the conventional tearing instability in Section 4. We discuss the application of the tearing instability in neutron stars and other astrophysical systems in Section 5. We conclude in Section 6.

In this work we have shown that the tearing mode instability operates under the Hall effect and resistivity in the Electron-MHD description. The appearance of the instability is similar to the usual MHD case, developing the characteristic reconnection islands, even though the mechanism is physically different, as the usual concepts of magnetic pressure and tension do not apply in this context. We find that the tearing instability facilitates a faster magnetic field decay, which is more evident for high Hall parameters, without leading to any significant amplification of the strength of the local magnetic field. Considering its role in neutron star magnetic field evolution, we have found it is more likely to occur just below the surface of strongly magnetised neutron stars or close to the crust-core boundary. In the first case the energetics of the instability are consistent with the amount of heat needed for a magnetar burst, which is likely to originate close to the surface, while the associated magnetic field strengths are sufficient to deform the crust. In the latter case, it may provide an extra channel for magnetic field decay and contribute to the quiescent emission. We note that the tearing instability discussed here may be relevant to other systems where evolution under the Hall effect and Electron MHD is important. Namely, the Hall effect is known to operate in protoplanetary discs \citep{Balbus:2001}. \cite{Lesur:2014} showed that the inclusion of ambipolar diffusion and Ohmic decay leads to the formation of magnetic zones and recently, \cite{Bethune:2016} showed that the magnetic field reverses direction within a narrow layer [c.f. Figure 7 of \cite{Bethune:2016}]. We speculate the these reversal regions may be appropriate sites for the development of the tearing instability with implications for the overall evolution of these protoplanetary discs. Observations of the magnetotail has provided evidence of reconnection activity in the region \citep{Nagai:2001, Runov:2003, Snekvik:2009} and the release of plasmoids due to the Hall effect \citep{Liu:2013}. While the system near the magnetotail is more complicated than the simple Electron-MHD evolution described here, the basic principles described here may be still in operation and enhance the reconnection and the subsequent plasmoid formation.

1607.07874

1607

1607.05040_arXiv.txt

We perform a detailed study of an effective field theory which includes the Standard Model particle content extended by a pair of Weyl fermionic SU(2)-doublets with opposite hypercharges. A discrete symmetry guarantees that a linear combination of the doublet components is stable and can act as a candidate particle for Dark Matter. The dark sector fermions interact with the Higgs and gauge bosons through renormalizable $d=4$ operators, and non-renormalizable $d=5$ operators that appear after integrating out extra degrees of freedom above the TeV scale. We study collider, cosmological and astrophysical probes for this effective theory of Dark Matter. We find that a WIMP with a mass nearby to the electroweak scale, and thus observable at LHC, is consistent with collider and astrophysical data only when fairly large magnetic dipole moment transition operators with the gauge bosons exist, together with moderate Yukawa interactions.

The conventional way to produce non-relativistic (cold) DM relic particle abundance, is the so called freeze-out mechanism~\cite{Lee:1977ua,Hut:1977zn}. Although this mechanism is well reviewed in the literature\cite{Hooper:2009zm, Jungman:1995df,Kolb:1990vq,Dodelson:2003ft,Weinberg:2008zzc,Lisanti:2016jxe}, it would be helpful to outline the main steps here. In the early universe, when the temperature was much higher than $M_D$, the would-be DM particles were in equilibrium, which means that it was equally possible to create and destroy pairs of them due to the $Z_2$-symmetry. As temperature of the universe was dropping, the thermal production of DM pairs became inefficient. Thus, $\chi_{1}^{0}$ pairs started to annihilate into lighter SM particles. As the number of these would-be DM particles was dropping, it became increasingly rare for them to interact with each other and annihilate. This yielded an almost constant number density of $\chi_{1}^{0}$ particles, which corresponds to DM relic density observed today. Assuming that $\chi_{1}^{0}$ is the lighter particle of the dark sector, one can evaluate the relic density accurately\footnote{Extensive discussion on the solution of the Boltzmann equation including coannihilation effects can be found in \cite{Edsjo:1997bg}.} by solving the corresponding Boltzmann equation: \begin{equation} \frac{d n_{\chi_{1}^{0}}}{dt}+3H n_{\chi_{1}^{0}}=-\vev{\sigma v_{rel}}\left(n_{\chi_{1}^{0}}^2 -n_{\chi_{1}^{0}}^{(eq)2} \right) \label{Boltz-eq}, \end{equation} where $H$ is the Hubble parameter defined as \begin{equation} H\equiv\dfrac{\dot{\alpha}(t)}{\alpha(t)}\;, \end{equation} and $\alpha(t)$ is the cosmic scale factor. Also $n_{\chi_{1}^{0}}$ is the WIMP number density and $n_{\chi_{1}^{0}}^{(eq)}$ is the corresponding quantity in equilibrium \begin{equation} n_{\chi_{1}^{0}}^{(eq)} \equiv g \left( \frac{m_{\chi_{1}^{0}} T}{2 \pi}\right)^{3/2} e^{-x}\;, \qquad x \equiv\frac{m_{\chi_{1}^{0}}}{T} , \label{neq} \end{equation} where $g$ is the number of the internal degrees of freedom of a particle, $\vev{\sigma v_{rel}}$ is the thermal average of the total annihilation cross-section of the WIMP to all allowed particles $(k,l)$, multiplied by the relative velocity of the incoming particles, which is usually expanded as \begin{equation} \vev{\sigma v_{rel}}=\displaystyle\sum_{k,l} \, \vev{\sigma_{\chi_{1}^{0}\chi_{1}^{0} \to k,l} v_{rel}} \ = \ a +b \, \vev{v_{rel}^{2}} \ + \ ... \label{csxv-non_rel_exp} \end{equation} It should be noted, that the second term on the r.h.s. of \eq{Boltz-eq} is responsible for creating $\chi_{1}^{0}$-pairs, while the first term for annihilating them. According to our description above, at high temperatures, much higher than $m_{\chi_1^0}$, the r.h.s of \eq{Boltz-eq} vanishes. This results to a constant particle number density since \begin{equation} \frac{d n_{\chi_{1}^{0}}}{dt}+3H n_{\chi_{1}^{0}} \ = \ \frac{1}{\alpha^{3}}\, \dfrac{d(\alpha^{3} n_{\chi_{1}^{0}})}{dt}=0 \;. \end{equation} For lower temperatures than $m_{\chi_1^0}$, the term $\vev{\sigma v_{rel}}n_{\chi_{1}^{0}}^{(eq)2}$ in \eq{Boltz-eq} should vanish, since the WIMP pairs are not produced effectively [see \eq{neq}]. Then the Boltzmann equation can be approximated as \begin{equation} \frac{d n_{\chi_{1}^{0}}}{dt} \approx -\left( \vev{\sigma v_{rel}}\, n_{\chi_{1}^{0}} +3H \right) n_{\chi_{1}^{0}} \;. \label{Boltz-eq2} \end{equation} The freeze-out temperature is defined as this where the annihilation rate becomes comparable to the expansion rate of the universe \begin{equation} \vev{\sigma v_{rel}}\, n_{\chi_{1}^{0}} \approx H \;. \end{equation} The freeze-out temperature $T_f$ can be evaluated iteratively, through \begin{equation} x_{f}=\log \left[ c(c + 2) \sqrt{\dfrac{45}{8}} \frac{m_{ \chi_{1}^{0} }M_{P} \left(a+6b/x_{f}\right)}{g^{1/2}_{\star} x_{f}^{1/2}} \right], \label{x_fo} \end{equation} where $x_{f}\equiv m_{\chi_1^0}/T_f$. The parameter $c$ is usually chosen $c \sim 0.5$, to get into agreement with precise numerical solutions of the Boltzmann equation. Furthermore, $M_{P} \approx 2.435 \times 10^{18} \; \GeV$ is the Planck scale, and $g_{\star}$ counts the relativistic degrees of freedom of the Standard Model at $T_{f}= {m_{ \chi_{1}^{0} }}/{x_{f}}$. It turns out that $x_{f} \simeq 25$. Calculating the freeze-out temperature, one can solve the Boltzmann equation and find the present WIMP relic density \begin{equation} \Omega h^2 \ \approx \ \frac{1.04 \times 10^9 \; \GeV ^{-1}}{M_{P}}\; \frac{x_{f}}{g^{1/2}_{\star}(a+3\, b \, x_{f}^{-1} )} \;. \label{relic-approx} \end{equation} For a WIMP mass at the electroweak scale, this formula becomes approximately $\Omega h^2 \approx 0.1 \frac{10^{-8} \; \mathrm{GeV}^{-2}}{a+3 \, b \, x_{f}^{-1}}$. From \eq{Planck} we get $\Omega h^2 \sim 0.1$, so the required cross-section is of order $ 10^{-8} \; \mathrm{GeV}^{-2}$ for $a=\mathcal{O}( \mathrm{GeV}^{-2})$, which is a typical EW cross section. If other particles are almost degenerate with WIMP, then there could be extra contributions (coannihilation effects) to the total annihilation cross-section due to them. Thus, the annihilation cross-section modified in order to incorporate these coannihilation effects~\cite{Griest:1990kh}. Following \cite{Griest:1990kh,Hooper:2009zm}, this change is \begin{equation} \displaystyle\sum_{k,l}\sigma_{\chi_{1}^{0}\chi_{1}^{0} \to k,l} \rightarrow \sigma_{eff}=\displaystyle\sum_{k,l} \; \displaystyle\sum_{i,j}\sigma_{i,j \to k,l} \; \dfrac{g_{i} g_{j}}{g^{2}_{eff}(x)}(1+\Delta_{i})^{3/2} \; (1+\Delta_{j})^{3/2} \; e^{-x (\Delta_{i}+\Delta_{j})} \;, \end{equation} where indices $i,j$ run over all the co-annihilating particles with $\Delta_{i}=\dfrac{m_{i}-m_{\chi_{1}^{0}}}{m_{\chi_{1}^{0}}} \lesssim 0.1$ and $g_{eff}(x)$ is defined as \begin{equation} g_{eff}(x) \equiv \displaystyle\sum_{i}g_{i} \, (1+\Delta_{i})^{3/2} \; e^{-x \Delta_{i}} \;. \end{equation} Such coannihilation effects, and other possible contributions to the relic abundance~\cite{Griest:1990kh}, have been included in our numerical analysis described in the following. \subsection{A close look at the relic density}\label{Close_Look-Section} Before discussing the bounds imposed by the data on $ \Omega h^2$, it would be helpful to study the numerical values of the annihilation cross-section that are used to calculate the relic abundance. As discussed in section \ref{sect:htogg}, if $M_D \gtrsim 90$ GeV, then the coupling to the Higgs boson is approximately zero. Therefore, the most important annihilation channels, assuming for the time being that coannihilation effects are irrelevant, are $\chi_{1}^{0} \chi_{1}^{0} \to W^{+}W^{-}$, $ZZ$, $ \gamma Z $ and $ \gamma \gamma$. There are no final states with fermions, since their corresponding interaction vertices are absent. There are no $\chi_{1}^{0} \chi_{1}^{0}Z / \gamma$ terms in the Lagrangian of \eqs{Zxx}{diag_dipole-3point}, or they are restricted because of bounds by direct detection experiments $Y^{h\chi_{1}^{0} \chi_{1}^{0}} \approx 0$. Keeping only the first term in the expansion of \eq{csxv-non_rel_exp} we obtain \begin{equation} a_{VV} = \frac{ \beta^{3/2}_{V}\, m_{\chi_{1}^{0}}^2 }{32\pi \, S_V\, v^2 }\, \frac{\left [g^2 v^4- 4\, g \, v^2 \omega \, K_{V} \left(m_{\chi_{1}^{0}} + m_{\chi} \right) +4\, K_{V}^2\, \omega^2 \left( 2m_{\chi_{1}^{0}} m_{\chi} +M^2_V \right) \right ]^2 }{v^6 \left( m_{\chi_{1}^{0}}^2 + m_{\chi}^2 -M^2_V \right)^2 }, \label{avv} \end{equation} where $V$ denotes $W$ and $Z$ gauge bosons in the final states for the processes $\chi_{1}^{0} \chi_{1}^{0} \to W^{+}W^{-}$ or $\chi_{1}^{0} \chi_{1}^{0} \to ZZ $. Also, we abbreviate, $\beta_V \equiv 1- {M_V^2}/{m_{\chi_{1}^{0}}^2}$, $K_{W} \equiv d_W$, $S_W \equiv 1$, $K_{Z}\equiv c_{W} \, (c_{W} d_W+s_{W} d_{\gamma})$ and $S_Z\equiv 2\,c_{W}^{4}$. The mass $m_{\chi}$ denotes $m_{\chi^{\pm}}$ for $V=W$ and $m_{\chi_{2}^{0}}$ for $V=Z$. For the channels $\gamma Z$ and $\gamma \gamma$, we find \begin{subequations} \begin{align} a_{\gamma Z} & = \frac{\beta_{\gamma Z}^3\, m_{\chi_{1}^{0}}^2 }{2 \pi\, c_{W}^2\, v^2} \frac{C_{\gamma}^2 \, \omega^2\, \left [ g\,v^2\left(m_{\chi_{1}^{0}} + m_{\chi_{2}^{0}} \right) - \omega\,K_{Z} \left( 4m_{\chi_{1}^{0}} m_{\chi_{2}^{0}} +M^2_Z \right) \right ]^2}{v^6 \left[ 2\left(m_{\chi_{1}^{0}}^2 + m_{\chi_{2}^{0}}^2\right) -M^2_Z \right]^2 } \, , \label{agammaz}\\ a_{\gamma \gamma} & =\frac{m_{\chi_{1}^{0}}^4\, m_{\chi_{2}^{0}}^2\, \omega^4\, C_{\gamma}^4}{\pi \,(m_{\chi_{1}^{0}}^2 + m_{\chi_{2}^{0}}^2)^2\, v^8} \, ,\label{agammagamma} \end{align} \end{subequations} with $\beta_{\gamma Z} \equiv 1- {M_Z^2}/{4m_{\chi_{1}^{0}}^2}$ and $C_{\gamma}\equiv (c_{W} d_{\gamma}-s_{W} d_W)$. These channels $\gamma\gamma$ and $\gamma Z$, contribute to the monochromatic gamma fluxes from the GC. Thus, in conjunction to the corresponding bounds from Fermi-LAT experiment, one gets severe constraints for the coupling $C_{\gamma}$. Due to absence of $\chi_{1}^{0}$ couplings to $Z$ and $\gamma$ and the nearly vanishing Higgs mediated $\hat{s}$-channel, all the above processes arise from $\hat{t}$ and $\hat{u}$ channels. \Eqsss{avv}{agammaz}{agammagamma}, contain one or more solutions with respect to $d_W$. This means that $d_W$ could act as a regulator that minimizes the total annihilation cross-section as the (required) low mass $M_D$ tends to amplify it (generally the cross section scales as $M_{D}^{-2}$ if we ignore magnetic dipole interactions). This minimization, will be proved essential when trying to obtain cosmologically acceptable relic abundance at the electroweak scale. Qualitatively, concerning the minimum of the {\it total} annihilation cross-section as a function of the dipole couplings one anticipates that each cross-section should be minimized for almost the same value of $d_W$, in order for the {\it total} annihilation cross-section to be at its minimum. In addition, $d_{\gamma}\approx \frac{s_{W}}{c_{W}} d_W$ so that $C_{\gamma}$ is quite small. This keeps $d_{\gamma}$ from obtaining large negative values, because $a_{WW}$ can be minimized only for $d_W>0$. \begin{figure}[t!] \centering \includegraphics[width=0.78\linewidth]{csxva-dw.png} \caption{\em The dependence of different annihilation channels on $d_W$ for $M_D =400 \GeV$, $\Lambda =1 \TeV$, $y=-y_{12}=-\frac{\xi_{12}}{2}=-0.8$ and $d_{\gamma}=0$. Notice that, in a certain range of $d_W$ values, there is at least one dip for each channel cross section.} \label{cross_sections_Vs_dw-a} \end{figure} \begin{figure}[H] \hspace*{-0.0cm} \begin{subfigure}[b]{0.52\textwidth} \includegraphics[width=\textwidth]{y12=-y_dwvsrelic.pdf} \caption{} \label{dwvsrelic} \end{subfigure} \begin{subfigure}[b]{0.52\textwidth} \includegraphics[width=\textwidth]{y12=-y_dgvsrelic.pdf} \caption{} \label{dgvsrelic} \end{subfigure} \vspace*{-1.5cm} \begin{center} \begin{subfigure}[b]{0.52\textwidth} \includegraphics[width=\textwidth]{y12=-y_MDvsrelic.pdf} \caption{} \label{MDvsrelic} \end{subfigure} \end{center} \caption{\em Relic abundance dependence on the parameters $(a) \, d_W$, $(b) \, d_{\gamma}$, $\text {and } (c) \, M_D$, for $\Lambda=1 \TeV$ and $y_{12}=-y$. The cosmologically allowed (shaded) region corresponds to the variation of the other parameters in \eqref{params} not shown in the plot. The horizontal line stands for $\Omega h^2 =0.12$.} \label{Omega_dependence} \end{figure} A numerical example is shown in Fig.~\ref{cross_sections_Vs_dw-a}. We observe that there are two minima for the annihilation cross-sections to $ZZ$, $W^{+}W^{-}$ and $\gamma Z$ and one minimum for $\gamma \gamma$. The first minimum of $a_{ZZ}$ and $a_{WW}$ coincides with the vanishing point of $C_{\gamma}$, which gives small cross-sections for $\chi_{1}^{0}\chi_{1}^{0} \to \gamma \gamma \text{ and } \gamma Z$. On the other hand, the second minimum of $a_{ZZ}$ and $a_{WW}$ is in a region where the annihilation to $\gamma\gamma$ and $\gamma Z$ blows up. Furthermore, for negative $d_W$, there are no such minima and, as can be seen from Fig.~\ref{cross_sections_Vs_dw-a}, every cross-section becomes quite large. Since \eq{relic-approx} is an approximation which could lead to an error up to $ \sim 10 \%$ (as discussed in \Ref{Griest:1990kh}), the Boltzmann equation must be solved numerically. To do this we implement the $d=4$ and $d=5$ operators to the computer program microOMEGAs\cite{Belanger:2014hqa} via the LanHEP\cite{Semenov:1998eb} package\footnote{More information about these packages can be found in https://lapth.cnrs.fr/micromegas/ and http://theory.sinp.msu.ru/$\sim$semenov/lanhep.html. } in order to obtain more accurate results for the relic abundance. In Figs.~\ref{dwvsrelic}, \ref{dgvsrelic} and \ref{MDvsrelic} we examine the dependence of the relic abundance $\Omega h^2$ on the parameters, $d_{W}$, $d_{\gamma}$ and $M_{D}$, respectively. Because all parameters in \eqref{params}, run freely, the corresponding plots are given as shaded areas in Fig.~\ref{Omega_dependence}. We remark that: {\it a)} The minimization effects on the various cross-sections discussed before, are evident in the numerical results too. {\it b)} As expected, when $M_D$ increases, $\Omega h^2$ increases too. {\it c)} For acceptable $\Omega h^2 $ and $M_D$ of a few hundred $\GeV$, $d_W$ must lie in the region $0.1 \lesssim d_W \lesssim 0.5$, which does not include the zero node. The dipole moment to photon $d_\gamma$ should be in the region $-0.2 \lesssim d_{\gamma} \lesssim 0.5$, which includes the zero node. {\it d)} The minimization of the total annihilation cross-section, is not enough to produce the observed DM density for $M_D \lesssim 200 \, \mathrm{GeV}$. \begin{figure}[H] \centering \includegraphics[width=0.70\linewidth]{y-dependence-ZZ.pdf} \caption{\em $ZZ$ annihilation cross section dependence on $d_W$ for different values of $y=-y_{12}=-\frac{\xi_{12}}{2}$, $\Lambda =1 \TeV$, $M_D =400 \GeV$, $d_{\gamma}=0$ and $v_{\mathrm{rel}}^2 =0.1$. At the minimum, the values of the cross-section decreases as we lower the values of $|y|$. The behaviour of the $W^{+}W^{-}$ annihilation channel is similar. } \label{csxvab-dw-MD=400_y} \end{figure} \begin{figure}[H] \hspace*{-0.3cm} \begin{subfigure}[b]{0.52\textwidth} \includegraphics[width=\textwidth]{y12=-y_y1vsrelic_MD=200GeV.pdf} \caption{} \label{yvsrelic_MD=200GeV} \end{subfigure} \begin{subfigure}[b]{0.52\textwidth} \includegraphics[width=\textwidth]{y12=-y_y1vsrelic_MD=400GeV.pdf} \caption{} \label{yvsrelic_MD=400GeV} \end{subfigure} \caption{\em $y$ vs. $\Omega h^2$, for $\Lambda =1 \; \TeV$ and $(a) \, M_D = 200 \; \GeV$ and $ (b) \, M_D = 400 \; \GeV$. Other parameters from the list \eqref{params} vary in the range constrained from ``Earth" constraints and for $y_{12}=-y$. The dependence of the relic density on $y$ changes for different values of $M_D$.} \label{Omega_dependence_y_MD=200|400} \end{figure} The dependence of the relic density on the parameter $y$ is complicated due to the following competing effects: The coannihilation channels, increase the total annihilation cross-section as $|y|$ tends to zero, since the mass differences of the initial particles involved become smaller and smaller. But, as shown in Fig.~\ref{csxvab-dw-MD=400_y}, the $b-$term in the expansion of \eq{csxv-non_rel_exp}, tends to decrease the value of the cross-sections (around the minimum), at least for the annihilation to $ZZ$ and $W^{+}W^{-}$. Moreover, in Fig.~\ref{Omega_dependence_y_MD=200|400} we study the dependence of $\Omega h^2$ on $y$, for various values of the mass $M_D$. In the region $M_D\lesssim 260 \;\GeV$, the relic abundance becomes smaller for smaller $y$ (an example for $M_D =200 \; \GeV$ is shown in Fig.\ref{yvsrelic_MD=200GeV}), which means that the coannihilation effects dominate over the $b-$term, and vice-versa for larger values of $M_D$ (Fig.\ref{yvsrelic_MD=400GeV}). \begin{figure}[t] \hspace*{-0.3cm} \begin{subfigure}[b]{0.52\textwidth} \includegraphics[width=\textwidth]{y12=-y_turning_point_y_MD=260GeV.pdf} \caption{} \label{turning_point_y} \end{subfigure} \begin{subfigure}[b]{0.52\textwidth} \includegraphics[width=\textwidth]{y12=-y_turning_point_dw_MD=260GeV.pdf} \caption{} \label{turning_point_dw} \end{subfigure} \caption{\em Turning point as can be seen in (a) $y$ versus $\Omega h^2$ and (b) $d_W$ versus $\Omega h^2$, for $y_{12}=-y$, $\Lambda =1 \TeV$ and $M_D = 260 \; \GeV$. The shaded area and the curves are as in Fig.~\ref{Omega_dependence}.} \label{turning_point} \end{figure} \begin{figure}[t] \centering \includegraphics[width=.52\linewidth]{y12=-y_y1vsrelic.pdf} \caption{\em $\Omega h^2$ versus $y$, for $y_{12}=-y$, $\Lambda =1 \; \TeV$ and $M_D \leq 500 \; \GeV$. The shaded area and the curves are as in Fig.~\ref{Omega_dependence}.} \label{Omega_dependence_y} \end{figure} There is a small region at $M_D \approx 260 \; \GeV$ where this dependence is mixed. We call this value of $M_D$ {\it ``turning point''}. An example of this behavior is shown in Fig.\ref{turning_point_y}. As we can see, the relic abundance rises until $y \sim -0.4$ and then decreases, but for $y\sim -0.06$ it starts to increase again. Also, as shown in Fig.\ref{turning_point_dw}, we obtain two maxima for $\Omega h^2$ with respect to $d_W$, as a result of this effect, since the value of $d_W$ which minimizes the annihilation cross-section depends on $y$. Although $y$ has no definite effect on $\Omega h^2$, the relic density increases as $M_D$ increases. Therefore, if we calculate the relic density in the allowed parameter space, the dependence of the relic on $y$ would be dominated by its dependence for larger $M_D$. In Fig.~\ref{Omega_dependence_y}, we show the dependence of $\Omega h^2$ on $y$. The relic density decreases as $|y|$ becomes larger and for $|y| \gtrsim 0.9$ the DM becomes under-abundant. Finally, the case where $y_{12}=0$ yields similar results with the case $y_{12}=-y$ just discussed, as can be deduced from Figs.~\ref{Omega_dependence} and \ref{Omega_dependence_y}. Also, for other values of the cut-off scale $\Lambda$, the parameters $d_{\gamma , \, W}$ and $y$ should be rescaled in order for the ratios ${d_W}/{\Lambda}$, ${d_{\gamma}}/{\Lambda}$ and ${y}/{\Lambda}$ to remain unchanged. \subsection{Cosmological constraints due to relic density} Having studied the constraints from $LEP$, $R_{h\to\gamma\gamma}$, the direct detection DM experiments as well as the Planck bound on the relic density for this effective theory, we are able delineate the cosmologically acceptable regions of the parameter space. For this reason, we perform a combined scan in the so far allowed parameter space which is also cosmologically preferred, for the cases $y_{12}=-y$ and $y_{12}=0$ at $\Lambda =1 \TeV$. \begin{figure}[t!] \centering\vspace{-.3cm} \includegraphics[width=0.52\linewidth]{y12=-y_dwvsdgamma_1TeV.pdf} \caption{\em The plane $d_W - d_{\gamma}$ of the parameter space that gives the observable relic abundance, for $\Lambda =1 \TeV$ and $y_{12}=-y$. The same region holds also for $y_{12}=0$. We allow variation of other parameters in \eqref{params} consistently with observational data.} \label{dw-dgamma} \end{figure} First, for $\Lambda =1\,\TeV$, in Fig.~\ref{dw-dgamma} we display the part of the $d_{\gamma} - d_{W}$ plane, that is compatible to the DM relic density, varying all the other parameters, but keeping $M_D \lesssim 500 \GeV$. Apparently the parameter $d_{W}$ is bounded to be positive in order to explain the DM relic abundance for a WIMP mass at electroweak scale. Also, the region where $d_{\gamma}$ is positive, is larger than the region where it is negative, a situation explained in the preceding analysis. A similar region is also found for $y_{12}=-y$ and $y_{12}=0$. \begin{figure}[t] \hspace*{-0.3cm} \begin{subfigure}[b]{0.52\textwidth} \includegraphics[width=\textwidth]{y12=-y_MDvsdw_1TeV.pdf} \caption{$M_D$ vs $d_W$, for $\Lambda =1 \; \TeV$ and $y_{12}=-y$.} \label{y12=-y_MDvsdw_1TeV} \end{subfigure} \begin{subfigure}[b]{0.52\textwidth} \includegraphics[width=\textwidth]{y12=0_MDvsdw_1TeV.pdf} \caption{$M_D$ vs $d_W$, for $\Lambda =1 \; \TeV$ and $y_{12}=0$.} \label{y12=0_MDvsdw_1TeV} \end{subfigure} \caption{\em As in Fig.~\ref{dw-dgamma} but for acceptable values on the plane $M_D - d_W$. } \label{MDvsdw} \end{figure} \begin{figure}[H] \hspace*{-0.3cm} \begin{subfigure}[b]{0.52\textwidth} \includegraphics[width=\textwidth]{y12=-y_MDvsy1_1TeV.pdf} \caption{$M_D$ vs $y$, for $\Lambda =1 \; TeV$ and $y_{12}=-y$.} \label{y12=-y_MDvsy1_1TeV} \end{subfigure} \begin{subfigure}[b]{0.52\textwidth} \includegraphics[width=\textwidth]{y12=0_MDvsy1_1TeV.pdf} \caption{$M_D$ vs $y$, for $\Lambda =1 \; \TeV$ and $y_{12}=0$.} \label{y12=0_MDvsy1_1TeV} \end{subfigure} \caption{\em Values on $M_D - y$ plane that provide acceptable DM relic abundance. } \label{MDvsy1} \end{figure} In Fig.~\ref{MDvsdw} we observe that $M_D$ vastly affects the allowed values for $d_W$ that provide the correct relic abundance. This is due to the fact that the minimum of the total annihilation cross-section depends on the mass $M_D$, as can be seen from \eqss{avv}{agammaz}{agammagamma} and also from the fact that the maximum of $\Omega h^2$ varies as $M_D$ changes, see also Fig.~\ref{MDvsrelic}. Moreover, as $M_D$ becomes larger, the minimization of the cross-section becomes less necessary. Note that, for $y_{12}=0$ there is a gap for $d_W$ at $M_D \approx 260 \GeV$, a result of the ``turning point'' discussed at the end of the previous paragraph (see Fig.~\ref{turning_point_dw}). For $y_{12}=-y$, this ``turning point'' is ineffective. \begin{figure}[t] \hspace*{-0.3cm} \begin{subfigure}[b]{0.52\textwidth} \includegraphics[width=\textwidth]{y12=-y_y1vsxi12_1TeV.pdf} \caption{$y$ vs $\xi_{12}$, for $\Lambda =1 \; \TeV$ and $y_{12}=-y$.} \label{y12=-y_y1vsxi12_1TeV} \end{subfigure} \begin{subfigure}[b]{0.52\textwidth} \includegraphics[width=\textwidth]{y12=0_y1vsxi12_1TeV.pdf} \caption{$y$ vs $\xi_{12}$, for $\Lambda =1 \; \TeV$ and $y_{12}=0$.} \label{y12=0_y1vsxi12_1TeV} \end{subfigure} \caption{\em As in Fig.~\ref{MDvsy1}, but for the Yukawa parameters $y-\xi_{12}$. } \label{yvsxi12} \end{figure} In Fig.~\ref{MDvsy1} one can see the dependence of $M_D$ on $y$, in the region where the DM density complies the current cosmological bound. We observe that for large values of $M_D$, for $|y|<~0.85 (1.25)$ for the case $y_{12}=-y \ (y_{12}=0)$ we obtain the desired $\Omega h^2$. On the contrary, when $M_D \lesssim 300 \GeV$ in both cases for $y_{12}$, $|y|$ seems to be strongly dependent on $M_D$. This happens because the bound on $|y|$ from Earth-based experiments becomes stronger than the one from the relic abundance for smaller masses. In addition to that, since $\Omega h^2$ tends to decrease as $|y|$ becomes smaller for $M_D \lesssim 260 \; \GeV$, $|y|$ is also bounded from below. Furthermore, due to the ``oscillation'' of the relic abundance (Fig.~\ref{turning_point_y}), at $M_D \sim 260 \; \GeV$ there is a ``gap'' on the allowed values of $y$ (similar to $d_W$). Additionally, in Fig.~\ref{yvsxi12}, we see that $\xi_{12}$ follows $y$, a remaining result from the direct detection bound (similar to Fig.~\ref{DirDet-region}). \begin{figure}[t] \hspace*{-0.3cm} \begin{subfigure}[b]{0.51\textwidth} \includegraphics[width=\textwidth]{y12=-y_masses_1TeV.pdf} \caption{} \label{} \end{subfigure} \begin{subfigure}[b]{0.50\textwidth} \includegraphics[width=\textwidth]{y12=-y_mass-diff_1TeV.pdf} \caption{} \label{} \end{subfigure} \caption{\em (a) The cosmologically allowed mass of the WIMP versus the mass of the heavy fermions and (b) their mass difference for $y_{12}=-y$. Similar regions can be obtained for $y_{12}=0$.} \label{mass-region} \end{figure} The Yukawa couplings and the mass parameter $M_D$ displayed, fix the masses and their differences. For the sake of completeness, the masses and their difference from $m_{\chi_{1}^{0}}$ are shown in Fig.~\ref{mass-region} for $y_{12}=-y$ (similar region holds also for $y_{12}=0$). We observe that $m_{\chi_{1}^{0}} \gtrsim 200 \; \GeV$, for $y_{12}=-y$, which is also what one should expect from Fig.~\ref{MDvsy1}. In addition to this, the mass difference $m_{\chi_{2}^{0}}-m_{\chi_{1}^{0}}$ is in the region $\sim 2 - 50 \; \GeV$. Finally, we note that this mass difference takes slightly larger values ($\sim 2-70 \; \GeV$ ) for the other case of the symmetric limit for $y_{12}$, while $m_{\chi^{\pm}}-m_{\chi_{1}^{0}}$ is always half that [see \eq{eq:spec}]. Accordingly, the smallest possible mass of the WIMP in this case is $\sim 250 \; \GeV$ (which again can be seen also from Fig.~\ref{MDvsy1}). \subsection{Gamma-rays}\label{Gamma-rays-sec} Having delineated the cosmologically acceptable regions concerning the DM abundance, we will proceed calculating other astrophysical observables, like the gamma-ray fluxes (monochromatic and continuous) originating from the Milky Way GC and dSphs. \subsubsection{Continuous Gamma spectrum}\label{Gamma-rays-Con} In our model the DM pair annihilation cross-sections have been studied in section \ref{Close_Look-Section}. In particular, the relevant relations can be found in \eq{avv}. From \Refs{Ackermann:2015zua, Hooper:2012sr} we observe that the bounds on the cross sections $a_{ZZ}$ and $a_{WW}$ are above the required $\sim 3 \times 10^{-26} \; \mathrm{cm}^{3} \mathrm{s}^{-1}$ (for masses above $200 \; \GeV$) which generally gives the desired relic abundance. More precisely, for $m_{\chi_{1}^{0}}\gtrsim 200 \; \GeV$, the bound from dSphs is below $\sim 5 \times 10^{-26} \; \mathrm{cm}^{3} \mathrm{s}^{-1}$ for the annihilation $\chi_{1}^{0} \chi_{1}^{0} \to W^{+}W^{+}$ (assuming that the branching ratio is $100 \%$). The same bound holds the annihilation to a pair of $Z$-bosons, since their gamma spectra are quite similar. When applied to our model, which generally gives smaller branching ratios, these bounds should be even weaker. \begin{figure}[t] \hspace*{-0.5cm} \begin{subfigure}[b]{0.53\textwidth} \includegraphics[width=\textwidth]{y12=-y_massvsdw_Contour_WW.pdf} \caption{ } \label{} \end{subfigure} \hspace*{.0cm} \begin{subfigure}[b]{0.53\textwidth} \includegraphics[width=\textwidth]{y12=-y_massvsCZ_Contour_ZZ.pdf} \caption{} \label{} \end{subfigure} \caption{ \em Allowed region in the parameter space from collider, DM direct detection and relic density constraints discussed in sections 4 and 5.3, respectively, as a function of the WIMP mass and the couplings $d_W$ and ${K_{Z}}/{c_{W}}$. The contours show the values of the thermally averaged cross-sections (a) for $a_{WW}$ and (b) for $a_{ZZ}$ in $\mathrm{cm}^{3}\mathrm{s}^{-1}$ for $y_{12}=-y$. Similar for $y_{12}=0$. We take $\Lambda =1\;\TeV$.} \label{WW-ZZ-gamma_rays_cont} \end{figure} As it is shown in Fig.~\ref{WW-ZZ-gamma_rays_cont}, the relevant to continuous emission of photons cross-sections, $\sigma_{\chi_{1}^{0}\chi_{1}^{0} \to W^{+}W^{-}, \, ZZ }$ are safe with experimental bounds from continuous gamma ray spectrum discussed in this paragraph. \subsubsection{Constraints from Gamma-ray monochromatic spectrum}\label{Gamma-rays-lines} As we have seen, this effective theory relies on the various WIMPs magnetic dipole moment operators in order to give us the observed relic abundance. This could result to annihilations of pairs of WIMPs into photons which could be detectable from observations of gamma ray monochromatic spectrum originated from the GC. In this paragraph, we will calculate the cross-sections for processes that could give such gamma rays (\eqs{agammaz}{agammagamma}). As input, we use the parameter space that evade all the other, previously examined, bounds and use the results from Fermi-LAT~\cite{Ackermann:2013uma,Ackermann:2015lka} to set additional bounds to the parameters of this model. \begin{figure}[H] \hspace*{-0.5cm} \begin{subfigure}[b]{0.53\textwidth} \includegraphics[width=\textwidth]{y12=-y_energyvsCgamma_Contour_gammagamma.pdf} \caption{} \end{subfigure} \hspace*{0.0cm} \begin{subfigure}[b]{0.53\textwidth} \includegraphics[width=\textwidth]{y12=-y_energyvsCgamma_Contour_gammaZ.pdf} \caption{} \end{subfigure} \caption{\em The allowed, as in Fig.~\ref{WW-ZZ-gamma_rays_cont}, region of the parameter space, in terms of the photon energy and the coupling $C_{\gamma}$. The contours show the values of the thermally averaged cross-sections $a_{\gamma \gamma}$ (a) and $a_{\gamma Z}$ (b) in $cm^{3}s^{-1}$ for $y_{12}=-y$. Again $y_{12}=0$ results in an almost identical plot.} \label{gammagamma-gammaZ-gamma_rays_lines} \end{figure} These bounds depend strongly on the DM halo profile\footnote{The bounds have up to a factor of $15$ difference for different profiles and regions.} (and the region of interest) that one follows. Thus, we study the profile which gives the strongest bound. This comes from the $R3$ region which is optimized for the Navarro-Frenk-White NFWc($\gamma=1.3$) profile~\cite{Navarro:1995iw} (the relevant discussion on these regions of interest is found in \cite{Ackermann:2013uma}). So, the annihilation cross-section for $\chi_{1}^{0}\chi_{1}^{0}\to \gamma \gamma$ for this region of interest is bounded to be smaller than $\sim 10^{-28} \; cm^{3}s^{-1}$ for photon energy ($E_{\gamma}=m_{\chi_{1}^{0}}$) at $200 \; \GeV$ up to $\sim 3.5 \times 10^{-28}$ for $E_{\gamma}\sim 450 \;\GeV$ (and if we extrapolate up to $\sim 5 \times 10^{-28}$ for $E_{\gamma}\sim 500 \;\GeV$). For the process $\chi_{1}^{0}\chi_{1}^{0}\to \gamma Z$, we need to rescale this bound by a factor of two, since there is one photon in the final state instead of two. This process results to different value of $E_{\gamma}=m_{\chi_{1}^{0}}\, (1- {m_{Z}^2}/{4m_{\chi_{1}^{0}}^2} )$. Fig.~\ref{gammagamma-gammaZ-gamma_rays_lines} illustrates that the annihilation to $\gamma Z$ (and less to $\gamma \gamma$), violates the Fermi-LAT bound, mainly for larger values of $E_{\gamma}$. Thus, the values of $d_{W}$ and $d_{\gamma}$ are constrained so that $C_{\gamma}$ is even smaller than the cosmologically acceptable values. \begin{figure}[t] \hspace*{-0.5cm} \begin{subfigure}[b]{0.53\textwidth} \includegraphics[width=\textwidth]{y12=-y_photon_couplingVSMD.pdf} \caption{} \label{MDvsCgamma} \end{subfigure} \hspace*{0.0cm} \begin{subfigure}[b]{0.50\textwidth} \includegraphics[width=\textwidth]{y12=-y_dwvsdgamma_1TeV_update.pdf} \caption{} \label{dwvsdgamma_update} \end{subfigure} \caption{\em Allowed regions on a) $M_D - C_{\gamma}$ plane and b) $d_W - d_{\gamma}$ plane for $y_{12}=-y$, consistent with ``Earth" constraints, the observed relic abundance and the bounds from gamma-ray monochromatic spectrum, discussed sections 4, 5.3 and 5.4.2, respectively, in the text. Almost identical regions are allowed for $y_{12}=0$. The contour lines in (b) show the value of the $\chi_{1}^{0}\chi_{2}^{0}$-photon coupling $C_{\gamma}$.} \label{fig:gammagamma} \end{figure} It is evident from Fig.~\ref{MDvsCgamma}, that in order this model to deceive the current monochromatic gamma ray bounds from GC, we should limit the dipole couplings so they satisfy the relation $|d_W s_{W}-c_{W} d_{\gamma}| \lesssim 0.05$ ($\Lambda =1 \; \TeV$) for $M_D = 200 \; \GeV$ up to $|d_W s_{W}-c_{W} d_{\gamma}| \lesssim 0.15$ for $M_D = 500 \; \GeV$. Therefore, one can delineate accordingly the parameter space on the $d_W - d_{{\gamma}}$ plane, that evades all bounds and yields the correct relic density, which is shown in Fig.~\ref{dwvsdgamma_update}. It should be noted, that the other parameters remain unchanged as in the previous section, since they do not affect WIMP pair annihilation rates to two photons or to a photon and a $Z$ boson. Other values of $y_{12}$ result to almost identical regions to these in Fig.~\ref{fig:gammagamma}. Concluding this paragraph, we note that the Fermi-LAT data set upper bounds to the annihilation cross-section of two WIMPs into one or two photons, relating strongly the two dipole couplings, resulting to positive values for $d_{\gamma}$. Therefore, the two neutral particles of the model have an almost zero coupling to photon ($C_{\gamma}\approx 0$), while the other parameters are intact. It is worth pointing out that there is a non-relativistic non-perturbative effect, known as ``Sommerfeld enhancement"~\cite{ANDP:ANDP19314030302}, that can boost the annihilation cross-section, sometimes even, by orders of magnitude. For the bi-doublet case here, it has been calculated in the literature and the results are shown in \Refs{Hisano:2003ec,Hisano:2004ds,Cirelli:2007xd}. As it turns out, for the masses we are considering here, this effect is non-important. It becomes only sizeable for WIMP, ``higgsino-like" masses greater than about $1~\mathrm{TeV}$ or so. \subsection{Neutrino flux from the Sun}\label{Neutrinos} Another interesting indirect signal could come from solar neutrino flux. The cross-section for neutrino production from WIMP annihilations in the Sun, can be decomposed to the spin-dependent and spin-independent WIMP-nucleon cross-sections. Therefore, such experiments compete with direct detection ones. Recent results from IceCube~\cite{Aartsen:2012kia}, show that the spin-independent cross-section bound is relaxed as compared to the one obtained from direct detection experiments~\cite{Akerib:2015rjg}. On the other hand, the latest spin-dependent cross-section bound from solar neutrino flux~\cite{Aartsen:2016exj}, is much stronger than the one derived from LUX~\cite{Akerib:2016lao} for $m_{\chi_{1}^{0}} \gtrsim 200 \GeV$. In our study, the spin-independent bound from IceCube is evaded, since the constraints from LUX have been introduced from the beginning of this analysis. In addition, due to the c.c. symmetry, the spin-dependent cross-section vanishes, since $\chi_{1}^{0 \dagger}\bar{\sigma}^{\mu}\chi_{1}^{0}Z_{\mu}$ is odd under the transformation introduced in section~\ref{sec:eft}. Thus, these bounds, leave the allowed parameter space unaffected.

\setcounter{equation}{0} \label{sec:concl} We have introduced in the SM particle spectrum a fermionic bi-doublet: a pair of Weyl fermion $SU(2)_L$-doublets, $\mathbf{D}_1$ and $\mathbf{D}_2$, with opposite hypercharges. In addition, we assume a discrete $Z_2$-symmetry that distinguishes $\mathbf{D}_1$ and $\mathbf{D}_2$ from the SM fields. This anomaly free set of fermions, together with the $Z_2$-symmetry are quite common features in non-supersymmetric SO(10) GUT constructions for light dark matter. Light $SU(2)_L$ doublets, whose components are parts of the WIMP have been also considered countless of times in ``UV-complete'' non-supersymmetric or supersymmetric models (\ie higgsino dark matter). Our work is related to these UV models when all other particles but the doublets have been integrated out in their low energy spectrum. At the renormalizable level the mass spectrum consists of a electromagnetically neutral, and a charged Dirac, fermions. Under the presence of $d=5$ operators, the neutral Dirac fermion is split into two Majorana states, the WIMP, $\chi_1^0$, and its excited state, $\chi_2^0$. Moreover, the $d=5$ operators include magnetic and electric dipole transitions which are, in principle, generated by a UV-complete theory, possibly at the TeV scale. We ask here the question whether the dark matter particle $\chi_1^0$, with mass $(m_{\chi_1^0})$, {\emph {around the EW scale}}, is compatible to various collider, astrophysical and cosmological data. In order to reduce fine tuning and extensive scans of the parameter space, in section~\ref{Spectrum-section} we adopted four scenarios, a,b,c and d, based on well motivated symmetry limits of the theory such as a charge conjugation or a custodial symmetry that act on $\mathbf{D}$'s and Higgs field $H$. These low energy symmetries simplify enough the analytical expressions of the interactions and possibly help to construct UV-completions of the model. After collecting all relevant $d=5$, and $d=6$ (though the latter not used in the analysis), operators in the \ref{sec:ope}, we went on to investigate their implications into collider and astrophysical processes. In section~\ref{sec:earth}, we performed a constraint analysis based {\it (i)} on scattering WIMP-nucleus recoiling experiments, such as LUX, {\it (ii)} on LEP searches for new fermions, as well as {\it (iii)} on LHC searches for the decay $h\to \gamma\gamma$. Bounds on the model parameters \eqref{params} are collected in Fig.~\ref{Combined}. Only in cases (b) and (c) there is still enough freedom to carry on. In the same section, we also studied contributions from the new fermion interactions into oblique electroweak $S,T$ and $U$ parameters. Only the $S$ parameter is affected, and, as a consequence, only case (a) is further constrained. Focused on the more interesting cases (b) and (c), in section~\ref{sec:astro} we calculated the relic density $\Omega h^2$ for $\chi_1^0$. In the presence of $d=5$ dipole operators there are destructive interference effects in the (dominant) amplitudes for WIMP annihilations (or co-annihilations) into SM vector bosons. The minima in the cross sections correspond to certain, usually non-zero, values for the coefficients of the dipole operators $d_W$ and $d_\gamma$ [see \eqs{agammaz}{agammagamma}]. Nearby these minima the relic density is found to be consistent with observation [\eq{Planck}] for $m_{\chi_1^0} \gtrsim 200 \GeV$. Although continuous gamma ray spectrum constraints are harmless, constraints from monochromatic gamma ray spectrum are serious for the photon dipole coupling as it is shown in Figs.~\ref{gammagamma-gammaZ-gamma_rays_lines} and \ref{fig:gammagamma}. The coefficient $d_W$ has to be more than $10\%$ a value which is non-negligible for UV models with Dark matter at the EW scale. $d_\gamma$ on the other hand can be tuned to zero without a problem. Apart from possible aesthetics, the main reason in insisting for EW dark matter mass, $m_{\chi_1^0} \approx M_Z$, has to do with enhancing the possibility of observing the dark sector at the LHC (or, in any case, to be as close as visible in the RunII phase). In section~\ref{sec:LHC} we estimated the cross section for producing $\chi_1^0$ at LHC with center of mass energy $\sqrt{\hat{s}} = 8, 13$ TeV and in association with a jet (monojet) or 2 jets or a $W$ or a $Z$. We found that the monojet process is the most promising with a few hundred of events at $\sqrt{\hat{s}} = 13 \TeV$ and with $m_{\chi_1^0} \simeq 200 - 350 \GeV$ (see Fig.~\ref{mono-jets}). Searching for dark matter and/or related particles at LHC consists in a major effort from physicists in high energy physics and astrophysics. An effective field theory for an electroweak dark matter described in this article may guide us closing that goal. \vspace*{3cm}

1607.05040

1607

1607.05089_arXiv.txt

{ We have carried out a high-precision astrometric analysis of two very-long-baseline-interferometry (VLBI) epochs of observation of the 13 extragalactic radio sources in the complete S5 polar cap sample. The VLBI epochs span a time baseline of ten years and enable us to achieve precisions in the proper motions of the source cores up to a few micro-arcseconds per year. The observations were performed at 14.4\,GHz and 43.1\,GHz, and enable us to estimate the frequency core-shifts in a subset of sources, for which the spectral-index distributions can be computed. We study the source-position stability by analysing the changes in the relative positions of fiducial source points (the jet cores) over a decade. We find motions of 0.1$-$0.9\,mas among close-by sources between the two epochs, which imply drifts in the jet cores of approximately a few tens of $\mu$as per year. These results have implications for the standard Active Galactic Nucleus (AGN) jet model (where the core locations are supposed to be stable in time). For one of our sources, 0615$+$820, the morphological and spectral properties in year 2010, as well as the relative astrometry between years 2000 and 2010, suggest the possibility of either a strong parsec-scale interaction of the AGN jet with the ISM, a gravitational lens with $\sim$1\,mas diameter, or a resolved massive binary black hole. }

Distant Active Galactic Nuclei (AGN), like quasars and BL Lacs, are currently used as position references in the definition of astronomical inertial frames, from radio \citep[the international celestial reference frame, ICRF][]{Fey,Fey2} to the optical (e.g., {\em Gaia}, \citealt{Linde}). The consolidation of reference frames at different regions of the spectrum relies on a well-defined and time-stable chromaticity (i.e., frequency dependence) of the AGN emission. It is well-known that the radio emission from AGN originates at relativistic jets with a frequency-dependent structure. The location of the peak intensity (frequently associated to the so called jet core) depends on the observing frequency due to synchrotron self-absorption effects \citep[e.g., ][]{Blandford,Lobanov}. Due to this effect, the position of the jet core gets closer to the AGN central engine as the observing frequency increases. The study of this self-absorption effect, also known as core-shift \citep[first found by][]{Marcaide89}, has important implications in the study of the jet physics \citep[e.g., ][]{Lobanov}, but it is also crucial for a proper panchromatic alignment of the AGN-based inertial reference frames \citep[e.g., ][]{Kovalev}. The sky location of AGN cores may not only depend on frequency, but also on time. If the opacity in the jet changes (owing to variability in the particle density and/or the magnetic-field structure) or the jet changes its orientation (e.g., owing to precession), the position of the core at any given frequency (and also the core-shift) evolves. This kind of an evolution of core positions encode information about the changing physical conditions at the innermost regions of the AGN jets, and also map into time-dependent misalignments among AGN-based reference frames at different frequencies. To date, a large fraction of geodetic and astrometric very-long-baseline-interferometry (VLBI) observations rely on the group-delay observable. The group-delay astrometry does not usually take the effect of source structures into consideration, whose time variability (and frequency dependence) can introduce astrometric biases of even several times the nominal astrometric group-delay precision \citep{Moor}. Restricting the observations to very compact jet structures \citep[i.e., jets with low ``structure indices''][]{Charlot} and/or to jets with smooth profiles in a particle-field energy equipartition, help us to minimize the frequency (and time) astrometry variations in the definition of the reference frames with group-delay astrometry \citep{Porcas}. But, in any case, the use of phase delays instead of group delays provides a better solution for accounting for the source structure in the astrometry. Moreover, the phase delays are more precise than the group delays by up to several orders of magnitude \citep[see][for a deeper discussion]{PaperIII,Tesis}. In recent decades, we carried out a set of very-long-baseline-array (VLBA) observations of the complete S5 polar cap sample \citep{Eckart} at 8.4, 15, and 43\,GHz \citep[][hereafter, Papers I and II, respectively]{PaperI, PaperII}. The S5 polar cap sample consists of 13 radio-loud AGN that are located at high declinations (circumpolar for the VLBA). The main goals of this campaign were the study of the frequency dependence and time stability of the jet structures (especially, the jet cores), as well as the characterization of the absolute kinematics of the optically-thin jet components of all sources. All observations were performed in phase-referencing mode, to enable us the use of differential phase-delays in the astrometry analysis of the source positions. The differential phase-delays are the most precise interferometric observables and encode robust information on the relative position of the sources \citep[e.g., ][]{Marcaide89}. We have published partial results about the evolving source structures at 8.4 and 15\,GHz (Paper I/II), as well as the first astrometry analysis at 15\,GHz, together with a description of our astrometry technique \citep[][hereafter Paper III]{PaperIII}. Astrometry results on small subsets of this source sample have also been reported \citep[e.g., ][]{P01,P00} This is the fourth paper in this publication series. In this paper, we report new results from the latest observations of this campaign, which were performed in year 2010 at two frequencies, 14.4 and 43.1\,GHz, using, for the first time in this project, the fast frequency-switching (FFS) observing capabilities of the VLBA \citep[see e.g., ][]{Middelberg}. In the next section, we describe our observations. In Sect. \ref{calibration}, we describe the calibration strategy. In Sects. \ref{results15} and \ref{results43}, we present our results at 14.4\,GHz and 43.1\,GHz, respectively. In Sect. \ref{results1543}, we compare the observations at both frequencies and present spectral-index images for a subset of sources. In Sect. \ref{conclusions}, we summarize our conclusions.

\label{conclusions} We report on quasi-simultaneous 14.4\,GHz and 43.1\,GHz VLBA observations of the S5 polar cap sample, performed in December 2010 in phase-referencing mode, using the fast-frequency-switching (FFS) capabilities of the VLBA, and compare them to earlier results at 15\,GHz band. We have performed a high-precision (differential phase-delay) analysis at 14.4\,GHz, solving for all the $2\pi$ phase ambiguities as in Paper III. Between the years 2000 and 2010, we find a $4.7\sigma$ proper motion of 42$\pm$9\,$\mu$as\,yr$^{-1}$ between the jet cores of sources 10 and 11. For the rest of source pairs, the separations did not change above 2\,$\sigma$. We have performed an SFPR calibration, from 14.4\,GHz to 43.1\,GHz, to determine the core shifts. Only nine of the thirteen sources could be imaged with this technique. We find typical core-shifts of 0.05$-$0.2\,mas. We have constructed robust spectral-index images of these nine sources. The spectral-index distributions follow the well-known steepening of the spectrum at the jet extensions, from an either flat- or inverted-spectrum regions associated to jet cores. There is one source, 0615$+$820, that shows a remarkable double structure at 43.1\,GHz (two components, one at northeast, NE, and one at southwest, SW), having one of them, NE, a prominent jet extension roughly perpendicular to the NE$-$SW direction. % Possible explanations for this intriguing source structure could be either a strong jet bending at parsec scales from the AGN central engine (due to interaction with the ISM), a gravitational lens with mas scale, or a binary massive black hole. The relative astrometry between NE and SW at 15\,GHz, using image over-resolution, shows a clear position drift of SW with respect to NE between years 2000 and 2010, thus supporting the third possibility (binary black hole). A deeper analysis of the results on this source, using all the available VLBI data, will be published elsewhere. Future observations at mm-wavelengths (with the Global mm-wave VLBI Array, GMVA) and at cm-wavelengths (using the RadioAstron satellite) are being planned.

1607.05089

1607

1607.03982_arXiv.txt

A galaxy group catalog is constructed from the 2MASS Redshift Survey (2MRS) with the use of a halo-based group finder. The halo mass associated with a group is estimated using a `GAP' method based on the luminosity of the central galaxy and its gap with other member galaxies. Tests using mock samples shows that this method is reliable, particularly for poor systems containing only a few members. On average 80\% of all the groups have completeness $>0.8$, and about 65\% of the groups have zero contamination. Halo masses are estimated with a typical uncertainty $\sim 0.35\,{\rm dex}$. The application of the group finder to the 2MRS gives 29,904 groups from a total of 43,246 galaxies at $z \leq 0.08$, with 5,286 groups having two or more members. Some basic properties of this group catalog is presented, and comparisons are made with other groups catalogs in overlap regions. With a depth to $z\sim 0.08$ and uniformly covering about 91\% of the whole sky, this group catalog provides a useful data base to study galaxies in the local cosmic web, and to reconstruct the mass distribution in the local Universe.

\label{sec:mass} One of the key steps in the halo-based group finder \citep{Yang2005a,Yang2007} is to have accurate estimates of the halo masses of candidate galaxy groups. As demonstrated in \citet{Yang2007}, halo mass is tightly correlated with the total luminosity of member galaxies. In practice, however, one can only estimate a characteristic luminosity which is the sum of the luminosities of member galaxies brighter than some given limit. For a relatively deep survey such as the SDSS, where the limit can be set sufficiently low, the characteristic luminosity is a good proxy of the total luminosity and so can be used to indicate halo mass. For a shallow survey like the 2MRS, on the other hand, only a few (in most cases one or two) brightest member galaxies in the halos can be observed. The characteristic luminosity is no longer the best halo mass estimator, and an alternative is needed. In this paper, we implement the `GAP' method proposed by \citet{Lu2015}. \subsection[]{The GAP halo mass estimator} \label{sec:gap} \begin{figure} \center \vspace{0.5cm} \includegraphics[height=8.0cm,width=9.0cm,angle=0]{f4.eps} \caption{$L_c-M_h(M_{\rm g})$ relation given by the MOCKg sample using abundance matching between the cumulative luminosity function of galaxies and the halo mass function. `Round 1' relation is obtained using all mock galaxies [blue dashed line, labelled as MOCKg(1)] while `Round 2' is obtained using central galaxies only [red line labelled as MOCKg(2)]. The true $L_c-M_h(M_{\rm t})$ relation given by the MOCKt sample is plotted with black solid points with error bars which indicate the $16\%-84\%$ percentiles of the distributions around the median values.} \label{fig:MhLcmock} \end{figure} \begin{figure*} \center \vspace{0.5cm} \includegraphics[height=10.5cm,width=11.5cm,angle=0]{f5.eps} \caption{Comparison between the original (open blue circles) $\log M_{\rm t}$ and corrected (solid circles) $\log M'_{\rm t}$ halo masses by using the luminosity gap between the central (brightest) and the second brightest (top left) and the fifth brightest member galaxies (top right), respectively. The error bars indicate the $16\%-84\%$ percentiles of the distributions. The standard variances $\sigma$ between estimated halo mass and the true halo mass are illustrated in the bottom panel. As the legend indicates, results are shown for groups with 2, 3, 4 and 5 members, while the halo masses estimated only by using central galaxies (member 1) are also given in the same panel. } \label{fig:M_correct} \end{figure*} \begin{deluxetable*}{lccccc} \tabletypesize{\scriptsize} \tablecaption{Parameters of the $\Delta \log M_g$ model obtained from mock 2MRS sample. [See Eqs. (\ref{eq:DM_func}) \& (\ref{eq:eta_abc})]} \tablewidth{0pt} \tablehead{$\Delta \log M_g$ & $\beta_1$ & $\alpha_2$ & $\beta_2$ & $\beta_3$ & $\gamma_3$ } \\ \startdata MEMBER 2 & $ 10.81^{+ 0.18}_{- 0.19}$ & $ 0.36^{+ 1.60}_{- 0.26}$ & $ -15.44^{ + 3.35}_{- 7.86}$ & $ 10.39^{+ 0.11}_{- 0.24}$ & $ 1.94^{+ 0.83}_{- 0.4 1}$ \\ \\ MEMBER 3 & $ 10.21^{+ 0.39}_{- 0.10}$ & $ 0.23^{+ 0.54}_{- 0.14}$ & $ -13.40^{ + 1.21}_{- 3.97}$ & $ 9.90^{+ 0.36}_{- 0.10}$ & $ 2.21^{+ 0.32}_{- 0.2 7}$ \\ \\ MEMBER 4 & $ 9.98^{+ 0.32}_{- 0.18}$ & $ 0.20^{+ 0.25}_{- 0.09}$ & $ -13.39^{ + 1.10}_{- 2.76}$ & $ 9.81^{+ 0.30}_{- 0.17}$ & $ 2.45^{+ 0.24}_{- 0.1 5}$ \\ \\ MEMBER 5 & $ 9.77^{+ 0.33}_{- 0.07}$ & $ 0.13^{+ 0.15}_{- 0.01}$ & $ -13.67^{ + 1.18}_{- 0.94}$ & $ 9.67^{+ 0.28}_{- 0.07}$ & $ 2.54^{+ 0.15}_{- 0.0 8}$ \\ \enddata \label{tab:DeltaMmock} \end{deluxetable*} In the `GAP' method, one first needs to estimate the $L_c$-$M_h$ relation. For MOCKt samples, since every groups have the true central galaxies and true halo masses from the simulation, we can obtain this relation directly. Hereafter we refer the $L_c-M_h$ relation obtained directly from the simulation as the intrinsic (true) relation. On the other hand, for observational samples, one can obtain this relation from the conditional luminosity function model \citep[e.g. ][]{Yang2003} or from halo abundance matching \citep[e.g.][]{Mo1999,ValeOstriker2006,Conroy2006, Behroozi2010, Guo2010}. Here we adopt the latter and assume that there is a monotonic relation between the luminosity of central galaxy and the mass of dark matter halo, so that a more luminous galaxy resides a more massive halo. We can then get an initial estimate of the dark matter halo mass for each central galaxy from \begin{equation}\label{eq:ab3} \int_{L_c}^\infty n_c (L_c') dL_c' = \int_{M_h}^\infty n_h (M_h') dM_h'\,, \end{equation} where, $n_c (L_c)$ is the number density of central galaxies with luminosity $L_c$ and $n_h (M_h)$ is the number density of halos (or halo mass function) with mass $M_h$. In this paper, we adopt theoretical halo mass function given by \citet{Tinker2008}. Note that, in this abundance matching approach, we also need to know whether a galaxy is a central or a satellite. Since we are trying to find galaxy groups within the observation (the 2MRS in our case), we can easily separate galaxies into centrals and satellites with the help of group memberships. As we will show later, although the $L_c-M_h$ relation we obtain may deviates from the true one, especially at the massive end, the deviation can be compensated to some extent by our `GAP'-based correction factor. Our modeling of the $L_c-M_h$ relation using Eq. (\ref{eq:ab3}) is carried out via the following two steps. First, before we are able to separate galaxies into centrals and satellites with the help of group memberships, we assume that all of them are centrals \citep[as shown in ][ more than 60\% of the galaxies are centrals]{Yang2008}. To show the performance, We have applied this to our mock 2MRS sample, and obtain the `Round 1' $L_c-M_{\rm g}$ relation, which is shown in Fig. \ref{fig:MhLcmock} as the blue dashed line. For comparison, we also plot, as black solid points, the true median $L_c-M_{\rm t}$ relation obtained from the true centrals and true halo masses in the simulation, with error bars indicate the $16\%-84\%$ percentiles of the distributions. Compared to the true relation, we see that, the Round 1 relationship shows a general agreement with the true one, with a slight over-prediction of the halo masses at the bright end and slight under-prediction at the faint end. The deviation at the massive end is caused by the Malmquist bias in the $L_c-M_{\rm g}$ relation which can be corrected by the `GAP' \citep[see ][]{Lu2015}. The deviation at the faint end is caused by the inclusion of all the galaxies (including satellites) in our abundance matching. As we apply our group finder to the galaxy catalog in the next step, the group membership will enable us to separate galaxies into centrals and satellites. We can then limit the application of the abundance matching to centrals only, and improve the $L_c-M_h$ relation. After two to three iterations we converge to a new set of group memberships and a new $L_c-M_{\rm g}$ relationship, which is referred to as 'Round 2' and shown as the solid red line in Fig. \ref{fig:MhLcmock}. After this step, there is no longer any systematic deviation of the $L_c-M_{\rm g}$ relationship relative to the true one at the low mass end. With the $L_c-M_{\rm g}$ obtained in this step, we can estimate the `luminosity gap', which is defined as the luminosity ratio between the central and a satellite galaxy in the same halo, $\log L_{\rm gap}= \log(L_c/L_s)$ \citep[see][]{Lu2015}. The halo mass is then estimated using the relation, \begin{equation}\label{eq:Mfunc} \log M_{\rm g}(L_c,L_{\rm gap}) = \log M_{\rm g}(L_c) + \Delta \log M_{\rm g}(L_c,L_{\rm gap}) \,. \end{equation} This halo mass estimator consists of two parts. The first part is an empirical relation between $M_{\rm g}$ and $L_c$ derived from Eq. (\ref{eq:ab3}) which is represented by the first term on the right side. Another part is the amount of correction to that relation, which is represented by the second term $\Delta \log M_{\rm g}(L_c,L_{\rm gap})$. In order to model this correction term, we use the following functional form, \begin{equation}\label{eq:DM_func} \Delta \log M_{\rm g}(L_c,L_{\rm gap}) = \eta_a \exp(\eta_b \log L_{\rm gap}) + \eta_c\,. \label{eq:delM} \end{equation} The parameters $\eta_a$, $\eta_b$ and $\eta_c$ all depend on $L_c$ as: \begin{eqnarray}\label{eq:eta_abc} \eta_a(L_c)~ &=&~ \exp(\log L_c-\beta_1) \nonumber \\ \eta_b(L_c)~ &=&~ \alpha_2(\log L_c +\beta_2) \\ \eta_c(L_c)~ &=&~ -(\log L_c - \beta_3)^{\gamma_3} \nonumber \end{eqnarray} which is specified by five free parameters. For a given $L_c - M_{\rm g}$ relation, we fit the model to the true halo masses $M_{\rm t}$ of our galaxy systems (groups) in our mock sample to have the minimum variances between $\log M_{\rm t}$ and $\log M_{\rm g}(L_c,L_{\rm gap})$ \citep[see][for details]{Lu2015}. Table \ref{tab:DeltaMmock} presents the set of best fit values of these parameters. Since the (mock) 2MRS sample is shallow, we provide the parameters up to 5 group members. As an illustration, Fig. \ref{fig:M_correct} shows the performance of this halo mass estimator. In the top two panels, the original $L_c-M_t$ relations are shown as the open circles; the GAP-corrected relations are shown as the solid points, with the left panel showing results for $L_s=L_2$ and the right for $L_s=L_5$. To see the improvement, we define a `pre-corrected' halo mass, \begin{equation}\label{eq:M'_h} \log M'_{\rm t} = \log M_{\rm t} - \Delta \log M_g(L_c,L_{\rm gap})\,. \end{equation} If the correction term $\Delta \log M_{\rm g}(L_c,L_{\rm gap})$ can perfectly describe the scatter in the original relation $L_c-M_{\rm t}$, then there would be no scatter in the $L_c-M'_{\rm t}$. We can see that, the scatter in the $L_c-M'_{\rm t}$ is significantly reduced compare to that in the $L_c-M_{\rm t}$ relation. For massive halos/groups, this improvement is quite notable where the scatter is reduced almost by a factor of two. The bottom panel of Fig. \ref{fig:M_correct} shows the standard deviation $\sigma$ of the halo mass $\log M_{\rm g}(L_c,L_{\rm gap})$ obtained by Eq.(\ref{eq:Mfunc}) from the true halo mass $\log M_{\rm t}$. In both \citet{Lu2015} and this paper, we find that using $L_5$ gives the best correction to the halo mass. As shown \citet{Lu2015}, such a correction factor is quite independent of the galaxy formation model used to construct the mock catalog. In this paper we use the set of best fit parameters only up to the fifth ranked member (see below). \subsection[]{The Group Finder} \label{sec:gf} The group finder adopted here is similar to that developed by \citet{Yang2005a}. It uses the general properties of dark matter haloes, namely size and velocity dispersion, to iteratively find galaxy groups. Tests show that this group finder is powerful in linking galaxies with dark matter halos, even in the case of single member groups. As we pointed out earlier, the halo mass estimation adopted in \citet{Yang2005a,Yang2007} is based on the ranking of the characteristic group luminosity and proves to be quite reliable for surveys like the 2dFGRS and SDSS. For the 2MRS considered here, we use the `GAP'-corrected estimator described above. The modified group finder with this halo mass estimator consists of the following main steps: ~~\\ \textbf{Step 1: Start the halo-based group finder.} \vspace{3mm} In the earlier version of the halo-based group finder, the first step is to use the FOF algorithm \citep{Davis1985} with very small linking lengths in redshift space to find potential groups. Here we assume all galaxies in our catalog are candidate groups. The halo mass of each candidate group is calculated using the $L_c-M_h$ relation obtained in Eq. (\ref{eq:ab3}) (Round 1). ~~\\ \textbf{Step 2: Update group memberships using halo information.} \vspace{3mm} After assigning halo masses to all the candidates, groups are sorted according to their halo masses. Starting from the most massive one, we estimate the size and velocity dispersion of the dark matter halo, using the halo mass currently assigned to it. A dark matter halo is defined to have an over-density of 180. For the WMAP9 cosmology adopted here, the radius is approximately \begin{equation} r_{180} = 1.33 \mpch \left( \frac{M_h}{10^{14}\msunh}\right)^{1/3} \left(1+z_{\rm group}\right)^{-1}, \end{equation} Here, $z_{\rm group}$ is the redshift of the group center. The line-of-sight velocity dispersion of the halo is \begin{equation}\label{eq:sigma} \sigma = 418\kms\left( \frac{M_h}{10^{14}\msunh}\right)^{0.3367}\,. \end{equation} Finally, following \citet[hereafter Y07]{Yang2007}, we use the luminosity weighted center of member galaxies as the new group center. Then, one can assign new member galaxies to the group according to the tentative group center, tentative estimates of halo size and velocity dispersion obtained in the above steps. The phase-space distribution of galaxies is assumed to follow that of dark matter, and the group center is assumed to coincide with the center of halo. We use the following function of the projected distance $R$ and $\Delta z = z - z_{\rm group}$ to describe the number density of galaxies at $z$ in the redshift space around the group center at redshift $z_{\rm group}$: \begin{equation} P_M(R,\Delta z) = {H_0\over c} {\Sigma(R)\over {\bar \rho}} p(\Delta z) \,, \end{equation} where $c$ is the speed of light and $\bar{\rho}$ is the average density of the Universe. We assume the projected surface density, $\Sigma(R)$, is given by a (spherical) NFW \citep{NavarroFrenkWhite1997} profile: \begin{equation} \Sigma(R)= 2~r_s~\bar{\delta}~\bar{\rho}~{f(R/r_s)}\,, \end{equation} where $r_s$ is the scale radius, and the shape function is \begin{equation} \label{fx} f(x) = \left\{ \begin{array}{lll} \frac{1}{x^{2}-1}\left(1-\frac{{\ln {\frac{1+\sqrt{1-x^2}}{x}}}}{\sqrt{1-x^{2}}}\right) & \mbox{if $x<1$} \\ \frac{1}{3} & \mbox{if $x=1$} \\ \frac{1}{x^{2}-1}\left(1-\frac{{\rm atan}\sqrt{x^2-1}}{\sqrt{x^{2}-1}}\right) & \mbox{if $x>1$} \end{array} \right.\,. \end{equation} The normalization of the profile depends on the concentration $c_{180}=r_{180}/r_s$ as: \begin{equation} \bar{\delta} = {180 \over 3} {c_{180}^3 \over {\rm ln}(1 + c_{180}) - c_{180}/(1+c_{180})} \,, \end{equation} where the concentration model of \citet{Zhao2009} is adopted. The redshift distribution of galaxies within the halo is assumed to have a normal distribution, and can be described as follows, \begin{equation} p(\Delta z)= {1 \over \sqrt{2\pi}} {c \over \sigma (1+z_{\rm group})} \exp \left [ \frac {-(c\Delta z)^2} {2\sigma^2(1+z_{\rm group})^2}\right ] \,, \end{equation} where $\sigma$ is the rest-frame velocity dispersion given by equation~(\ref{eq:sigma}). So defined, the three-dimensional density in redshift space is $P_M(R,\Delta z)$. Then, we apply the following procedures to assign a galaxy to a particular group. For each galaxy we loop over all groups, and compute the distances $R$ and $\Delta z$ between the galaxy and the group center. An appropriately chosen background level $B=10$ is applied to the density contrast for galaxies to be assigned to a group. If, according to this criterion, a galaxy can be assigned to more than one group it is only assigned to the one with the highest $P_M(R,\Delta z)$. Finally, if all members of two groups can be assigned to one, they are merged into a single group. Note that in our group finder, the background parameter $B=10$ is set to ensure the balance between the interlopers and completeness of group memberships. A lower $B$ value will increase both the completeness of the group memberships and the number of interlopers, especially in massive groups. Thus for those who care most about the completeness of the group memberships only, a lower value of $B$ (e.g. $B=5$) can be used \citep[see ][]{Yang2005a}. ~~\\ \textbf{Step 3: Update halo mass with `GAP' correction.} \vspace{3mm} Once the new membership to a group is obtained, we use the new central and satellite galaxy system to estimate the halo mass using the `GAP' method described by Eq. (\ref{eq:Mfunc}). For each candidate group, we use the $L_c-M_h$ relation and the luminosity gap $\log L_{\rm gap}$ between the central galaxy and the faintest satellite (if the group contains less than 5 members) or the fifth brightest galaxy (if the group has membership equal to or larger than 5), to estimate the halo mass. In practice, we only apply the luminosity gap correction for centrals in the luminosity range $10.5 \leq \log L_c \leq 11.7$. As shown in the top panels of Fig. \ref{fig:M_correct}, fainter ($\log L_c \leq 10.5$) central galaxies are basically isolated. For $\log L_c \geq 11.7$, we found that using the value of $L_c$ directly in the GAP leads to over-correlation. Thus, for these systems we set $\log L_c = 11.7$ to estimate the GAP correction. In addition, since our galaxy sample is magnitude limited to $K_s = 11.75$, our method also suffers from a `missing satellite' problem, in that some groups do not contain any satellites brighter than the magnitude limit. As an attempt to partly correct for this, we assume that each galaxy group that contains only one member (a central) has a potential member satellite galaxy with an apparent magnitude $K_s = 11.75$, which corresponds to a limiting luminosity $L_{\rm limit}$ at the distance of the group. A `GAP' correction, $\log L_c - \log L_{\rm limit}$, is also applied to all groups of single membership with $\log L_c - \log L_{\rm limit}\geq 0.5$, and the final halo mass of such a group is set to be the average value between this mass and the original mass based on the central galaxy alone. ~~\\ \textbf{Step 4: Update the $L_c - M_h$ relation and Iterate.} \vspace{3mm} Once all the groups have been updated for new memberships, we can distinguish between centrals and satellites. We use the updated central galaxy sample to update the $L_c-M_h$ relation (Round 2) to be used to assign halo masses to tentative groups. We iterate Steps 2-4 until convergence is reached. Typically three iterations are needed to achieve convergence. Our final catalog is the collection of all the converged groups with information about their positions, galaxy memberships, and halo masses. ~~\\ \textbf{Step 5: Update the final halo masses of groups.} \vspace{3mm} Once all the groups (memberships) have been finalized, we perform a final update of the halo masses of groups using an abundance matching method so that the halo mass function of the groups is consistent with theoretical predictions \citep[e.g.][]{Yang2007}. In performing the halo abundance matching, we measure the {\it cumulative} halo mass functions of groups following the procedures described in section \ref{sec:mock2MRS}.

In this paper, we have implemented, tested and applied a modified version of the halo-based group finder developed in \citet{Yang2005a, Yang2007} to extract galaxy groups from the 2MRS. Covering uniformly about 91\% of the sky, the 2MRS provides the best available representation of the structures in local universe, and so a group catalog constructed from it is useful for many purposes. However, 2MRS is quite shallow; in many cases only a few brightest members within a halo can be observed. To deal with this limit, we have updated the halo mass estimate used in the previous group finder with a new method based on `GAP'. This `GAP' estimator consists of two parts: (i) a relation between the luminosity of the central galaxy, $L_c$, and the halo mass, $M_h$, inferred iteratively from abundance matching between the luminosity of {\it central} galaxies and the masses of dark matter halos; (2) a luminosity gap correction factor obtained from the luminosity difference between the central galaxy and a faint satellite galaxy. In order to evaluate the performance of our modified group finder, we have constructed mock 2MRS galaxy samples based on the observed $K_s$-band luminosity function. The group catalog obtained from the mock 2MRS galaxy catalog shows a 100\% completeness for about $65\%$ of the most massive groups to $\sim85\%$ for groups with halo masses $\log M_h < 10^{14}\msunh$. On average, about 80\% of the groups have 80\% completeness. In terms of interlopers, about $65\%$ of the groups identified have none, and an additional 20\% have an interloper fraction lower than 50\%. Further tests on the halo mass estimation show that the deviation of the halo mass between the selected groups and the true halos is $\sim 0.35 {\rm dex}$ over the entire mass range. These tests demonstrate that the modified group finder is reliable for the 2MRS sample. Applying the modified halo-based group finder to the 2MRS, we have obtained a group catalog with a depth to $z \leq 0.08$ and covering $91\%$ of the whole sky. This 2MRS group catalog contains a total of 29,904 groups, among which 24,618 are singles and 5,286 have more than one member. Some of the basic properties of the group catalog are presented, including the distributions in richness, in redshift and in halo mass. This catalog provides a useful data base to study galaxies in different environments. In particular, it can be used to reconstruct the mass distribution in the local Universe, as we will do in a forthcoming paper.

1607.03982

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1607.06105_arXiv.txt

Large field models of inflation, in which the canonically normalized inflaton $\phi$ ranges over a distance $\geq m_{pl}$ in field space, have an appealing simplicity \cite{Linde:1983gd,Freese:1990rb}: they can be built from single field with a flat and even monomial potential, and some direct couplings to matter that the inflaton decays into at the end of inflation, reheating the universe. They generate the largest possible primordial tensor fluctuations, which could be detected by the new era of CMBR polarization experiments sensitive to gravitational radiation. The challenge such models present is that the infinite number of terms of the form \be \delta L \sim c_n {\phi^n}/{m_{pl}^{n-4}}\label{eq:annoying} \ee must have exquisitely small coefficients $c_n$ lest slow roll and vacuum energy dominance be spoiled. This is not a problem in perturbative quantum field theory, where such models can be radiatively stable, so long as $\phi$ is derivatively coupled to any heavy physics that it may reheat into. Even when graviton loops are included, models with small initial values of $c_n$ are still technically natural \cite{Linde:1987yb}, because an approximate shift symmetry \cite{Kaloper:2008fb,Kaloper:2011jz}\ becomes exact when $c_{n>0} = 0$. On the other hand, there is a body of theoretical evidence that quantum gravity does not allow for unbroken or weakly broken global symmetries, bringing into question the initial choice $c_n \ll 1$. Some mechanism which restricts the UV theory is required. Axion monodromy inflation \cite{Silverstein:2008sg,McAllister:2008hb,Kaloper:2008fb,Kaloper:2011jz}\ provides a candidate mechanism for UV completion of large field inflation. Most work on these models is by necessity rather technical, as it involves explicit string theory constructions in order to discuss the UV-complete theory in a controlled fashion. Our goal here is to understand the field theory description of what the successful constructions lead to. We will argue that monodromy inflation models use a hidden gauge symmetry which ensures the consistency of the low energy theory and protects it from large corrections descending from the ultraviolet (UV) completion. More specifically, the low-energy theory can be written as a massive $U(1)$ 4-form gauge theory \cite{Kaloper:2008fb,Kaloper:2011jz,Kaloper:2008qs,Marchesano:2014mla}, with the axionic inflaton dual to the longitudinal mode of the 4-form. Such massive gauge theories can arise from the old Julia-Toulouse description of phases which contains many on-shell topological defects \cite{Julia:1979ur,Quevedo:1996uu,Quevedo:1996tx}. We will review it and provide a new argument for the core ideas in \cite{Julia:1979ur,Quevedo:1996uu,Quevedo:1996tx}. The phases with condensed defects are not continuously connected to the usual perturbative vacuum of the original theory. The large inflaton field excursions encode {\it macroscopic} field configurations: large 4-form flux fluxes induced by a condensate of topological defects. We will show that this theory is well behaved at high energies, so that the UV completion enters only through irrelevant operators. The macroscopic nature of the field configurations ensures that these irrelevant operators do not spoil slow-roll inflation. The upshot is a London equation-level description of monodromy inflation, which complements the BCS-type constructions in \cite{Silverstein:2008sg,McAllister:2008hb,Kaloper:2008fb,Kaloper:2008qs,Marchesano:2014mla,DAmico:2012ji,DAmico:2012sz,DAmico:2013iaa}. The underlying effective theory is analogous to the massive gauge theory description of the ground state of a superconductor. This picture also sheds some light on the fate of the Weak Gravity Conjecture \cite{ArkaniHamed:2006dz}\ for massive gauge fields. We avoid the problems stemming from the sad fate of global shift symmetries in quantum gravity because there is no global shift symmetry to break. The discrete shift symmetry of the dual axion $\phi$ is not a remnant of a broken global symmetry, but a gauge symmetry which emerges from the dual gauge symmetry of the massive 4-form. As shown in \cite{Kaloper:2011jz}, this prohibits direct contributions of the form (\ref{eq:annoying}). One still needs to ensure that the construction of the massive U(1) 4-form gauge theory is consistent with the UV completion, but that is a more straightforward task, which should be addressed with the existing tools. While we focus here on generating quadratic axion potentials, we also discuss extensions which display flatter potentials, which emerge in the original string theory constructions \cite{Silverstein:2008sg,McAllister:2008hb,Dong:2010in,McAllister:2014mpa}. These can arise from coupling to additional fields controlled by sub-Planckian dynamical scales. With this motivation in mind, we will first discuss massive $U(1)$ vector fields, to highlight the essential physics in a more familiar context. In particular we will give a novel, and we feel compelling, motivation for the Julia-Toulouse description of the hydrodynamics of vortices.

CMBR observations will soon be able to efficiently constrain inflationary models at highest scales, or, with luck, discover them. This forces the question of the internal self-consistency of inflationary models. A UV complete model provides a proof in principle of such consistency, as well as a model with explicit, computable parameters. However it is clearly of interest if one can make general arguments as to which -- if any -- large field models are self-consistent, and why. Here we have pursued this question from the effective field theory point of view, and argued that massive abelian 4-forms provide robust effective field theories realizing axion monodromy inflation. Instead of a broken global shift symmetry, which is expected to be badly broken by quantum gravity, there is a nonlinearly realized gauge symmetry for the 3-form, and an associated discrete gauge symmetry for the dual inflaton. These control the possible UV corrections. We know of no principle in quantum gravity which prevents these gauge symmetries. Further, the large field excursions by the effective inflaton are duals of large gauge field fluxes. These large fluxes -- and in turn, the effective inflaton field values -- can be thought of as macroscopic quantities which characterize the size of the system, rather than high energy excitations of the inflaton. While there is an upper limit on just how large a flux can be, this would imply that it is controlled by the macroscopic properties of the low energy theory, rather than its UV features. The major question is whether the 3-form mass can be kept well below the Planck scale. In string theory models with high-scale supersymmetry, there will be $p$-form gauge symmetries that become linearly realized in the ultraviolet, which can become nonlinearly realized at a lower scale. The gauge mass and the presence of extra modes that may activate flattening of the inflaton potential are set by the string compactification and the IR dynamics that emerge. If the Julia-Toulouse mechanism \cite{Julia:1979ur,Quevedo:1996uu,Quevedo:1996tx}\ can be realized as IR dynamics in a UV-complete theory, this may provide an avenue to generate the needed scale hierarchies. Full string compactifications are one route to this. It would be interesting to find intermediate-scale four-dimensional models of Ginzburg-Landau type, in which one can see the transition from a massless to a massive gauge theory phase. Higher-dimensional intermediate models, such as the ``unwinding inflation" models of \cite{DAmico:2012ji,DAmico:2012sz}\ are also of great interest.\footnote{The $D3-{\bar D3}$ version in five dimensions has a clear relation to the present work. Here the D3-brane separation on the circle is the inflaton. If we T-dualize, the off-diagonal Wilson line dual to the D4 has a Chern-Simons coupling to the RR 4-form field strength, which upon compactification becomes the axion-4-form coupling we describe here.} \vskip.5cm {\bf Acknowledgments}: We would like to thank Guido D'Amico, Daniel Harlow, David E. Kaplan, Matthew Kleban, Emil Martinec, Hirosi Ooguri, Fernando Quevedo, Eva Silverstein and Lorenzo Sorbo for useful conversations on this and related subjects. N.K. thanks the CERN Theory Division and Particle Theory Group, University of Nottingham, for hospitality in the course of this work. Part of this work was carried out while A.L. attended the Aspen Center for Physics during the ``Primordial Physics" workshop, which is supported by National Science Foundation grant PHY-1066293. A.L. would like to thank the organizers, participants, and the ACP staff for a stimulating environment conducive to work. Part of this work was carried out while A.L. attended the ``Entanglement in Strongly-Correlated Quantum Matter" and ``Quantum Gravity Foundations: UV to IR" workshops at the KITP, during which his research was supported in part by the National Science Foundation under Grant No. NSF PHY11-25915. He would like to thank the organizers, participants, and KITP staff for a stimulating environment. N.K. is supported in part by the DOE Grant DE-SC0009999. A.L. is supported in part by DOE grant DE-SC0009987.

1607.06105

1607

1607.08610_arXiv.txt

In this article, we present a bouncing cosmology inspired by a family of regular black holes. This scale-dependent cosmology deviates from the cosmological principle by means of a scale factor which depends on the time and the radial coordinate as well. The model is isotropic but not perfectly homogeneous. That is, this cosmology describes a universe almost homogeneous only for large scales, such as our observable universe.

Singularities within the General Relativity (GR) context are still an open question in both cosmology \cite{Novello} and black holes (BH) physics \cite{Joshi}. A final answer could be the insertion of quantum effects to solve this problem. However, in several cases, it is possible to avoid the singularities with aid of violations in energy conditions. Such violations are more acceptable since the observation of cosmic acceleration \cite{Supernova,Supernova2}. By assuming these violations, the Hawking-Penrose theorems are not valid. Then, in both cosmology and BH physics, one may construct gravitational field solutions without a singularity at $t=0$, initial time, or $r=0$, the coordinate origin. The Bardeen \cite{Bardeen2} solution was the first regular black hole (RBH) developed. It is a compact object with an event horizon and without a physical singularity (see, for example, \cite{Ansoldi} and \cite{Lemos_Zanchin} for an introduction). This achievement is a consequence of ideas from Sakharov and co-authors \cite{Sakharov,Gliner} that the spacetime inside the horizon, where the matter has a high pressure, is de Sitter-like. That is, at the center there is a de-Sitter core which avoids the existence of a singular point. The Bardeen solution has spherical symmetry and it does not violate the weak energy condition (WEC). Later, Hayward \cite{Hayward} has constructed another regular solution with this symmetry. Such as the Bardeen RBH, it does not violate the WEC (recently \cite{Neves}, we present a spherical regular solution which violates the WEC). Currently, regular solutions with axial symmetry (BHs with rotation) have been created as well \cite{Various_axial,Various_axial2,Various_axial3,Various_axial4,Various_axial5,Neves2}. In our work \cite{Neves_Saa}, we have constructed a general class of regular solutions by using a mass function which depends on $r$, the rotation and a cosmological constant. In another front, physicists and cosmologists have developed several ways to avoid the initial singularity problem, i.e., the big bang. In GR context, in general, bouncing cosmologies violate some energy condition, such as the strong energy condition (SEC). According to these researches, sophisticated bouncing cosmologies\textemdash{}models beyond the standard $\Lambda$CDM model\textemdash{}manage to avoid the big bang problem and the standard model problems (flatness, homogeneity, horizon and isotropy problems), which are solved by adopting the inflationary mechanism by the $\Lambda$CDM model. In this list beyond the standard inflationary model, there exist, for example, the ekpyrotic cosmology \cite{Lehners} and the matter bounce model \cite{Brandenberger}. Moreover, according to Ref. \cite{Steinhardt}, several bouncing cosmologies without the inflationary mechanism do not suffer from the initial conditions and multiverse-unpredictability problems. In bouncing cosmology, attempts have been made to construct regular models by ignoring the homogeneity. And this assumption is appropriate from the observational point of view \cite{Wu}. That is, at large scales our universe is described as almost homogeneous, but the inhomogeneity increases at small scales (see Ref. \cite{Bolejko} for a review in inhomogeneous cosmology). For example, Ruiz and Senovilla \cite{Ruiz_Senovilla} show a class of bouncing or regular cosmologies by means of an inhomogeneous perfect fluid. These results generalize a previous work by Senovilla \cite{Senovilla}. Inspired by RBHs, in this work we are mainly interested to study the initial singularity problem by presenting a simple inhomogeneous and scale-dependent bouncing model. This simple model assumes a generalized scale factor similar to the mass function adopted in RBHs (with $r$-dependence) and a barotropic equation of state (EoS). The term in the metric which carries the scale-dependency may determine the regularity of this cosmology. Thus, the model may be regular for either a dust or a radiation dominated universes. However, the model may have a negative energy density at the bounce. According to \cite{Nemiroff}, this feature may be attributed to quantum effects when the universe was very small. In this sense, the energy conditions are not valid. The structure of this paper is as follows: in Section II we present some features of RBHs with spherical symmetry; in Section III we show the bouncing model with a $t$ and $r$-dependence in the scale factor; in Section IV the final remarks are presented. In this work, we have used the metric signature $diag(-+++)$ and the geometric units $G=c=1$, where $G$ is the gravitational constant and $c$ is the speed of light in vacuum.

1607.08610

1607

1607.01677_arXiv.txt

Astrophysical tests of the stability of dimensionless fundamental couplings, such as the fine-structure constant $\alpha$ and the proton-to-electron mass ratio $\mu$, are an optimal probe of new physics. There is a growing interest in these tests, following indications of possible spacetime variations at the few parts per million level. Here we make use of the latest astrophysical measurements, combined with background cosmological observations, to obtain improved constraints on Bekenstein-type models for the evolution of both couplings. These are arguably the simplest models allowing for $\alpha$ and $\mu$ variations, and are characterized by a single free dimensionless parameter, $\zeta$, describing the coupling of the underlying dynamical degree of freedom to the electromagnetic sector. In the former case we find that this parameter is constrained to be $|\zeta_\alpha|<4.8\times10^{-6}$ (improving previous constraints by a factor of 6), while in the latter (which we quantitatively compare to astrophysical measurements for the first time) we find $\zeta_\mu=(2.7\pm3.1)\times10^{-7}$; both of these are at the $99.7\%$ confidence level. For $\zeta_\alpha$ this constraint is about 20 times stronger than the one obtained from local Weak Equivalence Principle tests, while for $\zeta_\mu$ it is about 2 orders of magnitude weaker. We also discuss the improvements on these constraints to be expected from the forthcoming ESPRESSO and ELT-HIRES spectrographs, conservatively finding a factor around 5 for the former and around 50 for the latter.

Introduction} Now that we know (as a result of experiments at the Large Hadron Collider) that fundamental scalar fields are among Nature's building blocks \cite{ATLAS,CMS}, an obvious follow-up question is whether such scalar fields also play an role on cosmological scales. Among the astrophysical probes of such dynamical degrees of freedom, tests of the stability of Nature's dimensionless fundamental couplings are the most direct and model-independent. Whenever dynamical scalar fields are present, one naturally expects them to couple to the rest of the model, unless a yet-unknown symmetry suppresses these couplings \cite{Carroll}. In particular, a coupling to the electromagnetic sector will lead to spacetime variations of the fine-structure constant, which unavoidably imply a violation of the Einstein Equivalence Principle---see \cite{Uzan,GRG} for recent reviews. Astrophysical tests of the stability of fundamental couplings are an extremely active area of observational research. The deep conceptual importance of carrying out these tests has been complemented by recent (even if somewhat controversial) evidence for such a variation \cite{Webb}, coming from high-resolution optical/UV spectrosopic measurements of absorption systems along the line of sight of bright quasars. A significant effort is being put into independently confirming this result, including a dedicated Large Program with the Very Large Telescope's UVES spectrograph \cite{LP1,LP2,LP3}. Improving these tests is also a flagship science case for forthcoming facilities such as ESPRESSO \cite{ESPRESSO} and ELT-HIRES \cite{HIRES}; a roadmap for these future tests is discussed in \cite{GRG}. Meanwhile, an extensive range of theoretical models including varying couplings have been studied (see \cite{Uzan} for an overview). Among these is a class of phenomenological models first suggested by Bekenstein \cite{Bekenstein} where, by construction, the dynamical degree of freedom responsible for the varying coupling has a negligible effect on the cosmological dynamics. This includes the Sandvik-Barrow-Magueijo model for a varying fine-structure constant $\alpha$ (which stems from a varying electric charge) \cite{SBM} and the more recent Barrow-Magueijo model for a varying proton-to-electron mass ratio $\mu$ (which stems from a varying electron mass) \cite{BM}. The original works provided some qualitative constraints on the models \cite{SBM,BM}, an in particular kept the values of the cosmological parameters fixed at their best-fit values. This last assumption was also used in a more recent analysis of the $\alpha$ model \cite{Leal}. In this work we provide a more quantitative analysis, relying both on cosmological background data, astrophysical measurements of $\alpha$ and $\mu$ in the approximate redshift range $0<z<4.3$ and atomic clock laboratory bounds on the current drift rates of both quantities. These models are characterized by a single phenomenological dimensionless parameter, $\zeta$, describing the strength of the coupling of the dynamical scalar degree of freedom to the electromagnetic sector, and we therefore obtain improved constraints on this parameter for both models. This same parameter also determines the amount of Weak Equivalence Principle (WEP) violation in these models, and we also compare our astrophysical constraints on $\zeta$ to those that follow from local WEP tests. We note that in our analysis we will be considering the simplest version of both models. While extensions of the $\alpha$ model with additional (functional) degrees of freedom have been suggested \cite{extend1,extend2}, the quantity and quality of the available data (and the fact that no strong evidence for a nonzero coupling $\zeta$ currently exists) motivate us to restrict ourselves to the simplest scenarios. Still, future observations will certainly allow the testing of broader scenarios. Here we also provide an illustration of this, quantifying the improvements on constraints on the simplest $\alpha$ scenario expected from the ESPRESSO spectrograph.

1607.01677

1607

1607.01394_arXiv.txt

The stellar initial mass function (IMF) of early-type galaxies is the combination of the IMF of the stellar population formed in-situ and that of accreted stellar populations. Using as an observable the effective IMF $\aimf$, defined as the ratio between the true stellar mass of a galaxy and the stellar mass inferred assuming a Salpeter IMF, we present a theoretical model for its evolution as a result of dry mergers. We use a simple dry merger evolution model, based on cosmological $N$-body simulations, together with empirically motivated prescriptions for the IMF to make predictions for how the effective IMF of massive early-type galaxies changes from $z=2$ to $z=0$. We find that the IMF normalization of individual galaxies becomes lighter with time. At fixed velocity dispersion, $\aimf$ is predicted to be constant with redshift. Current constraints on the evolution of the IMF are in slight tension with this prediction, even though systematic uncertainties prevent a conclusive statement. The correlation of $\aimf$ with stellar mass becomes shallower with time, while the correlation between $\aimf$ and velocity dispersion is mostly preserved by dry mergers. We also find that dry mergers can mix the dependence of the IMF on stellar mass and velocity dispersion, making it challenging to infer, from $z=0$ observations of global galactic properties, what is the quantity that is originally coupled with the IMF.

\label{sect:intro} Understanding the properties and the origin of the stellar initial mass function (IMF) is currently one of the biggest challenges in galaxy formation theory. Observational constraints on the IMF provide us with a puzzling scenario. On the one hand the IMF appears to be remarkably self-similar across different environments within the Milky Way \citep[see e.g.][]{BCM10, Off15}. On the other hand, the IMF in early-type galaxies is inferred to vary systematically as a function of mass or velocity dispersion \citep{Tre++10,Cap++12,CvD12, Dut++12, TRN13, Spi++14, Pos++15}, although even for massive early-type galaxies we are still far from a clear picture \citep{SLC15}. Efforts have been put into reproducing from a theoretical standpoint the observed IMF trends. However, despite recent progress \citep{H+C11,Kru11,Hop12,GKH16}, we still lack a coherent description of star formation across all environments. One complication in comparing measurements of the IMF with models is that present-day stellar populations are ensembles of stars formed at different epochs in a range of environments. For massive early-type galaxies, a significant fraction of their present-day stellar mass is believed to be accreted from other systems \citep[e.g.][]{vDo++10}. If the IMF is not universal, then each accreted object will in general have a different IMF from the preexisting population of the central galaxy. The IMF of a massive galaxy at $z=0$ will then be the combination of the IMF of the stellar population formed in-situ and that of the accreted galaxies, possibly resulting in spatial gradients \citep{Mar++15, LaB++16}. How does this "effective" IMF evolve in time? Answering this question and comparing the predictions to observations provides a new way to test galaxy formation and star formation models. While the IMF is typically assumed as universal in cosmological simulations, there are studies of galaxy evolution based on semi-analytical models (SAMs) that allow for a non-universal IMF \citep{Nag++05, Bek13, Cha++15, GargiuloI++15, Fon++16}. These works explore mostly the effect of the IMF on the chemical evolution of galaxies. In order to isolate the effects of a varying IMF in the context of dry mergers, we adopt a simple model based on cosmological numerical simulations, galaxy-galaxy mergers simulations, and empirical prescriptions for the varying IMF. In practice, we use a simple prescription for assigning the starting ($z=2$) IMF of an ensemble of galaxies and then evolve the stellar population of central galaxies by merging their stellar content with that of accreted satellites. We tune our model to match the correlation between IMF normalization and stellar mass observed at $z\sim0$ and use it to make predictions on the stellar IMF of massive galaxies at higher redshifts. Though very simple in its construction, our model allows us to clearly isolate the effect of dry mergers, which are believed to represent the main growth mechanism of massive early-type galaxies at $z < 2$, on the evolution of the IMF. The paper is organized as follows. In \Sref{sect:model} we describe our model for the IMF of $z=2$ galaxies and its evolution as a result of dry mergers. In \Sref{sect:obs} we present low-$z$ measurements of the IMF used to calibrate our model. In \Sref{sect:results} we show our predictions on the time evolution of the IMF. We discuss our results in \Sref{sect:discuss}, while \Sref{sect:concl} concludes. We assume a flat $\Lambda$CDM cosmology with $H_0=70$~km~s$^{-1}$~Mpc$^{-1}$ and $\Omega_M = 0.3$. Throughout the paper velocity dispersions are expressed in units of km~s$^{-1}$.

\label{sect:concl} Using empirically motivated recipes for describing the IMF of massive galaxies together with the dry merger evolution model developed by \citet{Nip++12} we studied the time evolution of the effective IMF, defined as the ratio between the true stellar mass and the stellar mass one would infer assuming a Salpeter IMF, of a population of massive galaxies from $z=2$ to $z=0$. The models are set up to match the observed correlations between IMF and stellar mass and velocity dispersion in massive early-type galaxies at $z\sim0.3$. Our model predicts a decrease in the effective IMF of individual objects with time as galaxies merge with systems with a lighter IMF. This is seen in the evolutionary tracks in \Fref{fig:snap}. Trends between stellar mass and velocity dispersion are qualitatively preserved by dry mergers, but the slope of the correlation between stellar mass and effective IMF becomes less steep with time. At fixed velocity dispersion, the population of galaxies evolves with a constant or increasing IMF normalization at increasing $z$. This prediction appears to be in slight tension with the few observational constraints on the IMF available at $z>0.3$. However, there are systematic effects that could be affecting observations the imporance of which has not fully been assessed yet, most importantly the contribution of dark matter and radial gradients in the IMF. Assuming that the tension is real and not the effect of systematics, the most plausible way to reconcile model and observations is to allow the effective IMF of satellite galaxies to be drawn from a different distribution, heavier at fixed galaxy properties, than that describing centrals. Finally, we find that the relative dependence of the IMF on stellar mass and velocity dispersion gets mixed by dry mergers, making it difficult to observationally determine the fundamental parameter(s) governing the IMF from $z\sim0$ measurements.

1607.01394

1607

1607.04264_arXiv.txt

We report on the detection of a bright, short, structured X-ray burst coming from the supernova remnant RCW\,103 on 2016 June 22 caught by the \emph{Swift}/Burst Alert Telescope (BAT) monitor, and on the follow-up campaign made with \emph{Swift}/X-Ray Telescope, \emph{Swift}/UV/Optical Telescope and the optical/near infrared (NIR) Gamma-ray Burst Optical and Near-infrared Detector. The characteristics of this flash, such as duration and spectral shape, are consistent with typical short bursts observed from soft gamma repeaters. The BAT error circle at 68 per cent confidence range encloses the point-like X-ray source at the centre of the nebula, \oneEc. Its nature has been long debated due to a periodicity of 6.67 h in X-rays, which could indicate either an extremely slow pulsating neutron star, or the orbital period of a very compact X-ray binary system. We found that 20 min before the BAT trigger, the soft X-ray emission of \oneEc\ was a factor of $\sim$\,100 higher than measured 2 yr earlier, indicating that an outburst had already started. By comparing the spectral and timing characteristics of the source in the 2 yr before the outburst and after the BAT event, we find that, besides a change in luminosity and spectral shape, also the 6.67 h pulsed profile has significantly changed with a clear phase shift with respect to its low-flux profile. The UV/optical/NIR observations did not reveal any counterpart at the position of \oneEc. Based on these findings, we associate the BAT burst with \oneEc, we classify it as a magnetar, and pinpoint the 6.67 h periodicity as the magnetar spin period.

RCW\,103 is a shell supernova remnant (SNR) of $\sim$\,9 arcmin apparent diameter, expanding at around 1100 km s$^{-1}$, with an estimated age between 1350 and 3050 yr \citep{carter97}, and at a distance of 3.3 kpc \citep{caswell75}. \citet{frank15} gave a detailed, spatially resolved, account of the X-ray emission of the SNR, that can be modelled with an absorbed, non-equilibrium ionization state (NEI) plane shock model, with an average temperature of 0.58 keV and an absorbing column density $N_{\rm H}$\,=\,0.95\,$\times$\,10$^{22}$ cm$^{-2}$. The relative abundance of the most important metals (Ne, Mg, Si, S, and Fe) is generally found to be half the equivalent solar value, reflecting the stronger contribution of the metal-poor circumstellar medium emission with respect to the expected metal-rich ejecta. The compact soft X-ray source \oneEc\ (hereafter \oneE), which is the neutron star (NS) born from the core-collapse supernova explosion, lies nearly at the centre of the SNR. The association between the central compact object (CCO) and the SNR is proved by a depression of $\sim$\,1 arcmin in the H\textsc{i} emission of the SNR, which is positionally and kinematically coincident with the location of the CCO \citep{reynoso04}. However, the CCO has no confirmed counterpart at other wavelengths. The X-ray luminosity of \oneE\ can vary by more than one order of magnitude on a time-scale of years in the range 10$^{33}$--10$^{35}$ erg s$^{-1}$ \citep{deluca06}. The X-ray spectrum is rather soft, and it can be well described either by the sum of two blackbodies with temperatures of 0.5 keV (and corresponding emitting radius, $R_{\rm BB}$ of few hundred metres) and 1.0 keV ($R_{\rm BB}$ tens of metres), respectively, or by the sum of a soft black-body of 0.5 keV and a steep power-law of photon index $\sim$\,3 \citep{deluca06, esposito11}. One of its most enigmatic features is a periodicity of 6.67 h found in a long \emph{XMM-Newton} observation of the source \citep{deluca06}. It is debated if the periodicity refers to the rotational period of an extremely slow, and peculiar, NS, or it is an orbital modulation of an accreting compact X-ray binary system. In the first hypothesis the NS should have an extreme magnetic field ($B \sim$ 10$^{13}$--10$^{15}$ G) as is typical of the so-called magnetar systems. This could possibly explain the very long spin period because of a large spin-down due to the interaction with a fossil disc formed from the debris of the supernova (SN) explosion. In the latter case the system would be a quite odd example of a very young low-mass X-ray binary system, even if the requirement of an extreme magnetic field is probably still needed \citep{deluca06, li07, pizzolato08, bhadkamkar09, ikhsanov13}. In this paper, we report on the recent discovery of a bright X-ray flash observed on 2016 June 22 with the \emph{Swift}/Burst Alert Telescope (BAT) instrument \citep{gcn19547} from the RCW\,103 region. The position of the hard X-ray source responsible for the X-ray flash, labelled SGR 1617-5103, is compatible with the position of \oneE\ \citep{atel9180, atel9183}. We present a detailed spectral and timing study of the X-ray emission of \oneE\ before and after the BAT trigger, finding the CCO in an outburst state. We also report on the search for a transient UV/optical/NIR counterpart with the Gamma-Ray Optical and Near-infrared Detector \citep[GROND; ][]{Schady2016ATel9184} and \emph{Swift}/UV/Optical Telescope (UVOT) to the outbursting source, finding only upper limits at the position of the CCO. We propose the identification of SGR 1617--5103 with \oneE\ and we discuss the implications of this discovery for constraining the nature of \oneE.

We have reported on the characteristics of the \emph{Swift}/BAT bright X-ray burst coming from the direction of the RCW 103 nebula observed on 2016 June 22. At the same time, we have also presented spectral and timing results concerning the X-ray activation of the CCO at the centre of the RCW~103 nebula, \oneE, and the rapid follow-up in NIR, optical, and UV bands in our search for a possible counterpart at these wavelengths. The long-term X-ray history of this source has already established flux variations of about two orders of magnitude on a years-long time-scale \citep{gotthelf99, deluca06}. The last outburst of this source happened between 1999 and 2001, when the source reached an intensity possibly similar to the one observed here \citep{garmire00}, and then began a slow return to its pre-outburst luminosity on a years-long time-scale \citep[see fig.~2 in ][]{deluca06}. This time we had the chance to closely monitor the start of the outburst and its short-term evolution, because \emph{Swift} was triggered by the detection of an X-ray burst in the direction of \oneE. As the characteristics of this burst (duration, spectral shape, and total fluence) are typical of soft gamma repeaters (SGR; whereas for the same reasons an association with a type-I X-ray burst is ruled out), and, at the same time, \oneE\ showed a dramatic change in flux, spectral shape and folded profile, we shall consider \oneE~ as the source originating the burst detected by BAT and take this as evidence for associating this peculiar CCO with the class of the magnetars \citep{duncan92, thompson93, thompson95}. We note that a very similar line of evidence was sufficient to grant the magnetar status to a relatively low ($B \sim$\,5\,$\times$\,10$^{13}$ G) magnetic field pulsar, PSR J1846--0258, in Kes 75 \citep{gavriil08}. This discovery makes the small group of CCO objects rather inhomogeneous based on the values of their inferred magnetic fields. CCO sources with $B$-field estimates show rather low values \citep[$B$\,$\lesssim$\,10$^{11}$ G; ][]{halpern10,gotthelf13} compared to typical values found in young NSs in high-mass binaries. \citet{gotthelf08} coined the term of \emph{anti-magnetars}, in antithesis to the supercritical $B$-field values of magnetars, to designate the CCOs hosted in SNRs. Although \oneE\ was not considered among the CCO sources listed in \citet{halpern10} because of its soft X-ray variability, it still fulfils all the other criteria for a CCO classification. It is most probably a \textit{classical high $B$-field ($B$\,=\,10$^{14}$--10$^{15}$ G) )} magnetar \citep{deluca06}, and, even considering a scenario where the initial spin-down was driven by an ejector phase of \textit{magnetized} debris, the required dipole field would still be above 10$^{12}$ G \citep{ikhsanov13}. This suggests that it can be difficult to generalize and assume all CCO objects as young and very low magnetised NSs \citep[see e.g.,][for magnetars hosted in a SNR environment]{gaensler01, vink08,gao16}, unless very ad hoc criteria are chosen. We studied the soft X-ray evolution of the source in the first 3 weeks of the outburst thanks to the monitoring campaign of \emph{Swift}. The X-ray light curve shows a clear peak just close to the time of the BAT trigger, with a steep decrease in the following hours, until a plateau is reached within 1 d from the X-ray burst. The flux evolution in the following weeks did not show any evident sign of fading, suggesting, as in the previous outburst, a possible slow decay to the pre-outburst luminosity levels. The initial steep decay, and the flatter evolution is similar to what observed in the case of other magnetars outbursts \citep[see e.g. the decay of the SGR 1E 2259+58 in its 2002 June outburst; ][]{woods04}, and, more generally, closely resembles the behaviour of transient magnetars like the SGR 1627--41, that shows similar flux variations on similar time-scales \citep{esposito08}. We tracked the most significant spectral changes using as a benchmark the time-averaged \textit{quiescent} spectrum of \oneE\ from the \emph{Swift}/XRT observations taken about 2 yr prior to 2016 June 22. We did not choose to disentangle the \oneE\ emission from the contribution of its nebula, because of the intrinsic bias and dependence of the results from the choice of the source and background extraction regions. Instead we modelled both components in a single fit to the data, using the results from the extensive work on the SNR emission made by \citet{frank15}. In this way, we obtained a statistically acceptable description of the data, and we could constrain at much higher confidence the parameters determining the spectral state of the source. The \oneE\ emission along all the outburst showed little variation in the time-averaged spectral shape, characterized by a soft thermal component of temperature $\sim$\,0.6 keV and a hard X-ray tail, carrying about 10\% of the total source emission. The spectral shape in the first observation of June 22, 20 min earlier than the BAT burst, seems to be harder than the other late-time spectra, but we also note that observations performed a few hours after the BAT event showed a rapid return to the temperature of $\sim$\,0.6 keV that also characterized the pre-outburst spectrum. The outlined spectral characteristics such as peak thermal temperature, harder flux excess during the outburst, time-scale of flux variations are all in agreement with the general properties shown by transient magnetars \citetext{see e.g.\, \citealp{kaspi03,scholz11,scholz12}, or these general reviews: \citealp{rea11, mereghetti15, turolla15}}. We studied the timing characteristics of the pulsed profile of \oneE, comparing the profiles at different times. We remark that, given the short time-span of the observations during the outburst, we were not able to clearly detect the 6.67 h periodicity, however, it is evident that any reasonable change in its value cannot have any statistically significant effect on the folded profile. A sort of bimodality in the pulsed profile was shown by the sparse observations of this source in the 1999--2005 years, where it was already found that when \oneE\ was in a brighter state the profile was remarkably different, and more structured \citep{deluca06}. We have observed that this change is not gradual, but it happens on a very short time-scale at the time of the outburst peak (Fig.~\ref{fig:folding}). The folded profile in outburst clearly shows that a significant phase change took place, and similarly to other magnetars where the same phenomenology is present \citep{kaspi03,woods04,dib09,woods11}, it could indicate a general re-arrangement of the magnetic field, which also caused the rapid dissipation of energy in the burst event. Clearly, it is comparatively much more difficult to assess if the sudden change in the folded profile is also associated with a frequency glitch, as observed in many magnetars \citep{dib08}, as the fractional frequency shifts are generally less than a part in a million except in some exceptional cases \citep{palmer02}. Future observations, spanning a longer time-frame, will hopefully set a constraint on this issue. Because of this characteristic change in the pulsed profile, commonly observed after a burst in magnetars \citep{mereghetti15}, we believe that the 6.675 h periodicity cannot be of orbital origin as had being speculated in \citet{bhadkamkar09}, but it must be associated with the NS spin period. Early suggestions for the presence of $dips$ in the folded profile \citep{becker02} are ruled out, as the spectral hardening appears strongly correlated with the total flux over the entire flux range, and it is not localized in the bottoms of the folded profile (Fig.~\ref{fig:hardness}). This finding makes \oneE\ the slowest pulsar to our knowledge \citep[the second being RX J0146.9+6121, with a spin period of 1380 s; ][]{haberl98}, and also makes \oneE\ a rather exceptional object among all the known magnetars, because the distribution of spin periods of these objects lies in only a decade of periods between $\sim$\,2 and 12 s \citep{olausen14}. Because of this extreme slow spin, the present rotational energy stored in the NS would be $\sim$ 3.4\,$\times$\,10$^{37}$ erg (assuming a canonical moment of inertia 10$^{45}$ g cm$^{2}$), and this value is very close to the energy dissipated in the burst event ($\sim$\,2\,$\times$\,10$^{37}$ erg), thus indicating that magnetic dissipation of an intense field must be the main, if not the only source responsible for the observed burst radiation. If the spin period of the NS at its birth was similar to that of other magnetars, some mechanism must have furiously braked it down on a time-scale comparable with the SNR age, which is only $\sim$ 2000 yr. Within the magnetar scenario, \citet{deluca06} proposed that the braking could be provided by a propeller effect due to a reservoir of mass formed from the SNR material fallback \citep[a fossil disc, see also][]{wang06}, continuously expelled at the magnetospheric radius of the NS. The only constraint which appears reasonable is that the NS initial period should have been longer than 0.3 s to avoid the disc disruption by the relativistic outflow of the newly born active radio pulsar. A magnetar with a magnetic field $B$\,=\,5\,$\times$\,10$^{15}$ G could reach the period observed in \oneE, after having expelled 3\,$\times$\,10$^{-5}$ M$_{\odot}$, in less than the age of the SNR. It is interesting to note that \citet{reynoso04} found the presence of an H\textsc{i} depression region around \oneE\ of radius 64 arcsec, and a lack of evidence for a possible ionized H\textsc{ii} region. The missing mass was evaluated to be $\sim$\,0.3 M$_{\odot}$, thus suggesting that a strong sweeping of material at the centre of the SNR might have taken place. \citet{li07} has further explored this scenario through a Monte Carlo simulation of a population of 10$^{6}$ NS magnetars interacting with a fallback disc. The NS population differs in initial spin periods, axis orientations, $B$-field, and mass of the fallback disc. He found that most of the magnetars ($\sim$\,99 per cent) would be found 2500 yr after their birth in the ejector phase (when the radiative pressure from the NS keeps the surrounding plasma away from the light cylinder and the spin-down can only be provided by magnetic dipole emission), but that 0.6\% could be found in the propeller phase (when the disc radius is between the magnetospheric and the light cylinder radius) and be effectively braked to periods $>$\,10$^{3}$ s as possibly happened for \oneE. Alternatively, the \oneE\ could be a binary system formed by a very low mass star and a magnetar with a spin (quasi-)synchronous with the orbital period \citep{pizzolato08}. In this scenario the torque needed to slow down the NS is provided by the interactions between the magnetic field and the surrounding material, similar to the case of white dwarfs in intermediate polars. However, the presence of a low-mass companion which could have survived the supernova explosion, is rather unlikely, given that such mass-lopsided system would have been prone to unbinding. In this context, it is relevant that \citet{deluca08} ruled out all but late M-type stars as possible companions, while \citet{li07} showed that, even in the case of survival, an irradiation-induced wind would not be able to power the observed X-ray emission. In search of possible counterparts at other wavelengths, we have also shown the results from a rapid optical/NIR follow-up of \oneE\ with GROND made 0.26 h after the BAT trigger and by \emph{Swift}/UVOT in the optical/UV bands both at the time of the trigger and in the whole set of \emph{Swift} observations of the source. No significant counterpart was detected in the stacked images consistent with the position of \oneE. We derived a series of upper limits on the magnitudes in the different bands from NIR to UV, thus supporting the absence of any irradiated close companion star, or distant accretion flow \citep{wang07}.

1607.04264

1607

1607.02442_arXiv.txt

This is a preliminary report on surface photometry of the major fraction of known globular clusters, to see which of them show the signs of a collapsed core. We also explore some diversionary mathematics and recreational tables.

A focal problem today in the dynamics of globular clusters is core collapse. It has been predicted by theory for decades \citep{hen61,lyn68,spi85}, but observation has been less alert to the phenomenon. For many years the central brightness peak in M15 \citep{kin75,new78} seemed a unique anomaly. Then \citet{aur82} suggested a central peak in \object{NGC 6397}, and a limited photographic survey of ours \citep[Paper I]{djo84} found three more cases, \objectname{NGC 6624}, \objectname[M 15]{NGC 7078}, and \object[Cl 1938-341]{Terzan 8}), whose sharp center had often been remarked on \citep{can78}. As an example of how the new AASTeX object tagging macros work, we will cite some of the ``Superlative'' objects mentioned in section 10 of Trimble's (1992) review of astrophysics in the year 1991. The youngest star yet found was \object[\[JCC87\] IRAS 4]{IRAS 4} in \objectname{NGC 1333}. \object{70 Oph} was found to be the longest period spectroscopic binary. The most massive white dwarf was \object{GD 50}, estimated at 1.2 solar masses. The first neutral hydrogen found in a globular cluster was \object{NGC 2808} while the \objectname[SDSS J093401.92+551427.9]{I Zw 18} retained the record for metal deficiency. However, another low metallicitity galaxy was \object{UGC 4483} in the \objectname{M 83} group. The largest redshift \object[PC 1247+3406]{source} in 1991 was found at z=4.897. Lastly, what paper would be complete without a mention of the \object[M1]{Crab nebula}!

1607.02442

1607

1607.05324_arXiv.txt

We study the behaviour of the dynamical and stellar mass inside the effective radius ($r_{e}$) of early-type galaxies (ETGs). We use several samples of ETGs -ranging from 19 000 to 98 000 objects- from the ninth data release of the Sloan Digital Sky Survey. We consider Newtonian dynamics, different light profiles and different Initial Mass Functions (IMF) to calculate the dynamical and stellar mass. We assume that any difference between these two masses is due to dark matter and/or a non Universal IMF. The main results for galaxies in the redshift range $0.0024 < z < 0.3500$ and in the dynamical mass range 9.5 $<$ log(M) $<$ 12.5 are: i) A significant part of the intrinsic dispersion of the distribution of dynamical vs. stellar mass is due to redshift. ii) The difference between dynamical and stellar mass increases as a function of dynamical mass and decreases as a function of redshift. iii) The difference between dynamical and stellar mass goes from approximately 0\% to 70\% of the dynamical mass depending on mass and redshift. iv) These differences could be due to dark matter or a non Universal IMF or a combination of both. v) The amount of dark matter inside ETGs would be equal to or less than the difference between dynamical and stellar mass depending on the impact of the IMF on the stellar mass estimation. vi) The previous results go in the same direction of some results of the Fundamental Plane (FP) found in the literature in the sense that they could be interpreted as an increase of dark matter along the FP and a dependence of the FP on redshift.

\label{sec:intro} The measurement of masses of galaxies has been, over a long period of time, an interesting and difficult problem, which has elicited the application of various and diverse techniques \citep{spi75,bur75,sof01,sim07}. Since the determination of rotation curves for a large number of spiral galaxies \citep{sof01} and the suggestion that these rotation curves are flat because of the presence of an unseen amount of mass which has been called `dark matter', the determination of the mass of all types of galaxies has become a pressing concern of modern Astronomy. It is fair to say at this point that there is no direct evidence of the existence of dark matter and that there are other explanations which, although not as currently popular as dark matter, may explain the observations quite reasonably. The total mass of a galaxy is composed of two elements; luminous matter and dark matter. If we assume that both luminous and dark matter respond to the Newtonian gravitational law in the same way, then the difference between the dynamical mass and the luminous mass of a galaxy provides us with an estimation of the amount of dark matter present in the galactic system in question. From such a determination we would be able to study if a dependence of the amount of dark matter with dynamical mass and/or redshift exists. Measuring the amount of radiation from a particular galaxy, combined with typical mass to light ratios (${\bf M}$/$L$) that have been calibrated using different stellar samples in our own Galaxy, allows us to estimate its stellar, gas and dust content. Moreover, rotation curves for spiral galaxies permit the calculation of dynamical mass inside any radius for which a value of rotation velocity is known, allowing us, in principle, to calculate from these two determinations the amount of dark matter present in the galaxy under study. As is well known, rotation velocity curves are used for studying the kinematics of galaxies, determining the amount and distribution of mass interior to a given radius, to derive an insight into galactic evolutionary histories and the possible role that interactions with other systems may have played. Since rotation curves may be obtained at different wavelengths they provide information as to the kinematics of different constituents of a galaxy. They may be observed in the infrared as well as in the optical, which may be used to trace ionised gas and the stellar motions, also in the radio and microwave regimes which trace the neutral and molecular gas components of a galaxy. Recently, stellar population synthesis models have been used to calculate galactic masses. These models also give us an idea of the total stellar content of a galaxy as well as the distribution of stars of all the different spectral types and luminosity classes \citep{aug10,bar11,son12}. Dynamical theoretical models can also be used to calculate masses for early-type galaxies (ETGs), such as those which \citet{vandermarel1991} constructed for 37 bright elliptical galaxies. From these models he found an average (M/$L)_B=(5.93 \pm 0.25)h_{50}$. Discrepancies of the observed velocities in the outer parts with those predicted by the models may be explained by the inclusion of massive dark haloes. \citet{gerhard2001} performed dynamical studies of the shapes of line-profiles for 21 elliptical galaxies; they used them to investigate the dark halo properties and dynamical family relations of these galaxies. They appear to have minimal haloes implied from the fact that the ratio M/$L_B$ turned out maximal. Some of these galaxies showed no dark matter within $2r_e$. \citet{cap06} investigated the correlations between the mass-to-light M/$L$ ratios of 25 elliptical and lenticular galaxies. Field and cluster galaxies presented no difference, and their dark matter content within an effective radius $r_e$ was $\sim 30\%$ of the total mass contained there. It appeared that the amount of dark matter correlates with galactic rotation velocity; in the sense that more massive slow-rotating galaxies contain less dark matter that the fast-rotating galaxies. There have been many papers in which dynamical arguments are used to calculate the dynamical mass of galaxies, and hence, by comparison with the amount of luminous mass, they calculate the amount of dark matter present, see for example: \citet{thomas2007}, \citet{williams2009}, \citet{thomas2011}, \citet{cap13} to mention a few. Also check the detailed introduction to this topic published in \citet{nig15}. The gravitational lens phenomenon provides direct and precise measurements of masses of galaxies at different scales, and allows us to establish the nature and presence of dark matter in a galactic system. Elliptical galaxies have been considered to have extended dark-matter massive haloes \citep{tre10} that follow the \citet{nav96} density profiles. \citet{ber93} and \citet{hum06} have studied the kinematics of different components in nearby elliptical systems and have concluded that dark matter haloes are required to explain the dynamics of massive elliptical galaxies, provided that Newtonian gravity be valid at these scales. Galactic mass determinations have also been made using weak and strong lensing observations \citep{lag10,hoe05,gav07}. The fraction of total mass in the form of dark matter in ETGs, $f_{DM}$, appears to increase with growing radius reaching values of $\sim 70\%$ at five effective radii \citep{tre04}. Furthermore, $f_{DM}$ within a fixed radius seems to grow with galaxy stellar mass and with velocity dispersion \citep{tor09,nap10,gra10,aug10a}. $f_{DM}$ varies from small values as in the case of bright giant elliptical galaxies \citep{rom03} to very large values, as has been found for dwarf spheroidal galaxies by \citet{sim07}. Studies of the Virgo giant elliptical galaxy NGC 4949 (M60) by \citet{teo11} reveal that the kinematics of Planetary Nebulae in this object is consistent with the presence of a dark matter halo with $f_{DM} \sim 0.5$ for $r=3r_e$. \citet{deb01} presented three-integral axisymmetric models for NGC 4649 and NGC 7097 and concluded that the kinematic data for NGC 4649 only require a small amount of dark matter, however \citet{das11} determine $f_{DM} \sim 0.78$ at $r=4r_e$ for NGC 4649. Using gravitational lensing experiments, \citet{koopmans2006} find a projected dark matter fraction of $<f_{DM}>=0.25 \pm 0.06$ for 15 ETGs, while \citet{bar11} studying sixteen early-type lens galaxies determine the lower limit for dark matter $f_{DM}$ inside the effective radius. The median value for this fraction is $12 \%$ with variations from almost 0 to up to $50 \%$. As mentioned above, direct detection of dark matter has not been achieved yet. Its presence requires the validity of Newtonian gravity. If we were to assume that at these very low acceleration regimes Newtonian gravity is not valid or may be slightly modified \citep{mil83} then further developments have explained several phenomena without the need of dark matter e.g. spiral galaxies, flat-rotation curves \citep{san02}, projected surface density profiles and observational parameters of the local dwarf spheroidal galaxies \citep{her10,mcg10,kro10}, the relative velocity of wide binaries in the solar neighbourhood \citep{her12}, fully self-consistent equilibrium models for NGC 4649 \citep{jim13} and references within among others. In this paper we present a study of luminous and dynamical mass inside the effective radius of ETGs considering Newtonian dynamics. We search for differences between these masses and assume that any difference is due to dark matter or a non-Universal IMF or a combination of both. The structure of this study is as follows; in \S 2 we present the sample of ETGs used in this work, in \S 3 we discuss the calculation of the stellar and virial masses for the galaxies in the sample, in \S 4 we discuss the distribution of stellar mass as a function of virial mass, in \S 5 and \S 6 we outline the difference between virial and stellar mass as a function of mass and redshift, in \S 7 we discuss our results in the Fundamental Plane context and finally in \S 8 we present the conclusions.

The analysis of the distribution of stellar mass with respect to virial mass on several samples of ETGs from the SDSS DR9 in the redshift range $0.0024 < z < 0.3500$ and in the dynamical mass range $9.5 < log({\bf M_{virial}}/{\bf M_{\odot}}) < 12.5$ has yielded the following results: \begin{itemize} \item A significant part of the intrinsic dispersion of the distribution of log(${\bf M_{virial}/{\bf M_{\odot}}})$ vs. log$({\bf M_{*}/{\bf M_{\odot}}}$) is due to redshift (see Fig. 4). \item The difference between dynamical and stellar mass depends on mass and redshift. \item The difference between dynamical and stellar mass increases as a function of dynamical mass and decreases as a function of redshift. \item The difference between dynamical and stellar mass goes from almost zero to approximately 70\% of the virial mass depending on mass and redshift (see figure 5). This difference is due to dark matter or a non-universal IMF or a combination of both. \item The amount of dark matter inside ETGs would be equal to or less than the difference between dynamical and stellar mass depending on the impact of the IMF on the stellar mass estimation. \end{itemize} The previous results have been analysed in the FP context and we have found that they go in the same direction as some FP results found in the literature in the sense that they could be interpreted as an increase of dark matter along the FP and a dependence of the FP on redshift. However in this work we have considered that the dark matter follows the same density profile as the stellar component and that the scale factor K in the virial relation (see equation 1) depends only on the light profile which, according to some authors \citep{cio96,koopmans2006,thomas2011}, is not appropriate for massive and compact galaxies respectively. In a forthcoming paper we will analyse these variables and their possible relation with the log(${\bf M_{virial}/{\bf M_{\odot}}})$-log$({\bf M_{*}/{\bf M_{\odot}}}$) difference.

1607.05324

1607

1607.03114_arXiv.txt

During their formation phase stars gain most of their mass in violent episodic accretion events, such as observed in FU\,Orionis (FUor) and EXor stars. V346\,Normae is a well-studied FUor that underwent a strong outburst beginning in $\sim$\,1980. Here, we report photometric and spectroscopic observations which show that the visual/near-infrared brightness has decreased dramatically between the 1990s and 2010 (${\Delta}R\approx10.9^{\rm m}$, ${\Delta}J\approx7.8^{\rm m}$, ${\Delta}K\approx5.8^{\rm m}$). The spectral properties of this fading event cannot be explained with variable extinction alone, but indicate a drop in accretion rate by 2-3 orders of magnitude, marking the first time that a member of the FUor class has been observed to switch to a very low accretion phase. Remarkably, in the last few years (2011-2015) V346\,Nor has brightened again at all near-infrared wavelengths, indicating the onset of a new outburst event. The observed behaviour might be consistent with the {\it clustered luminosity bursts} that have been predicted by recent gravitational instability and fragmentation models for the early stages of protostellar evolution. Given V346\,Nor's unique characteristics (concerning outburst duration, repetition frequency, and spectroscopic diagnostics), our results also highlight the need for revisiting the FUor/EXor classification scheme.

FU\,Orionis stars (FUors) are pre-main-sequence stars that undergo optical outbursts, interpreted as extremely active phases of mass accretion ($\sim10^{-4}\,$M$_{\sun}$\,yr$^{-1}$). These sources are classified by the observation of outbursts that increase the optical/infrared brightness by up to 6\,mag for several decades as well as their spectroscopic characteristics, which include CO overtone bandhead $2.3$\,$\mu$m absorption and P\,Cygni H$\alpha$~profiles that point towards strong outflow activity \citep[for a recent review see][]{aud14}. Only about a dozen stars are generally considered to belong to the FUor class, with a few more dozen FUor candidates. Besides their extremely high accretion luminosities, FUor outbursts are characterized by their very long outburst durations \citep[estimated to $\sim10^2$...$10^3$\,years;][]{har96} - in fact the commonly accepted FUor members are still in outburst since their discovery decades ago, including the prototype FU\,Orionis itself. The long outburst durations are also used to separate FUors from the EXors, which are eruptive stars of shorter outburst decay time (months to few years) and with lower accretion rates \citep[$\sim10^{-6}\,$M$_{\sun}$\,yr$^{-1}$,][]{her89}. V346\,Normae is in the Sa~187 cloud at a distance of $700$\,pc and is generally considered to be a genuine member of the FUor class \citep{pru93,rei97a}. Outburst activity of V346~Nor was first discovered in 1983 when \citet{gra85} reported the appearance of a star-like object in the visible ($V=16$\,mag). \citet{abr04} compiled the light curve and concluded that the near-infrared brightness was increasing continuously throughout the 1990s. The FUor classification is also based on the detection of strong Li\,I absorption, a P Cygni-like H$\alpha$ profile, and the detection of strong water vapour absorption around 1.9\,$\mu$m \citep{rei85}. Here, we report photometric observations (Sect.~\ref{sec:observations}) which show that V346\,Nor has experienced a dramatic drop in visual/infrared excess, indicating that the object has switched from a high-accretion to a low-accretion phase. We discuss our observed dramatic variability in Sect.~\ref{sec:results} and close with a summary and discussion on the broader implications of our study in Sect.~\ref{sec:conclusions}.

\label{sec:conclusions} Our study shows that V346\,Nor has switched from an extreme FUor-type outburst into a much less active phase, consistent with a change in the accretion rate by 2-3 orders of magnitude. With ${\Delta}R=10.9$\,mag and ${\Delta}K=5.8$\,mag, the observed dimming event far exceeded the variability that has been observed on other FUors before \citep[e.g.\ ${\Delta}R\approx3$\,mag, V899\,Mon;][]{nin15}. Therefore, V346\,Nor opens for the first time the opportunity to study an FUor system in post-outburst and to investigate how the disc adjusted to the sudden drop in accretion luminosity. For instance, some valuable insights might arise from modelling the observed anti-correlated variability at near-infrared/mid-infrared wavelengths. Also, the short-wavelength flux (R-band) started to decay a few year ahead of the near-infrared flux (e.g.\ J/H/K-band; Fig.~\ref{fig:photometry}, top), which might indicate that the disc needed a few years to react to the sudden drop in accretion luminosity. With an outburst duration of 20...25\,years, the V346\,Nor outburst was shorter than expected for FUors ($10^{2}$...$10^{3}$\,years), but significantly longer than for EXor (up to a few years). This suggests that eruptive stars might not be well-represented by the classical bimodal FUor/EXor classification scheme, but exhibit a more continuous range of properties. Already a few other objects with ``intermediate'' outburst durations of a few years have been identified \citep{aud14,con16} and started to blur the conventional distinctions between the FUor/EXor class. However, V346\,Nor represents the most extreme case, in particular as the object entered already into a new outburst, less than a decade after the decay of its FUor-type eruption. This recurrence cycle is orders-of-magnitudes shorter than the typical time span of $\sim$5000...50,000\, years that has been proposed for classical, isolated FUor outbursts \citep{sch13}, but might be consistent with the {\it clustered luminosity bursts} that have been predicted by recent disc gravitational instability and fragmentation models for the earliest stages of protostellar evolution \citep[e.g.][]{vor15}. Besides its atypical lightcurve, V346\,Nor is also ambiguous in its spectral diagnostics, as it features Br$\gamma$ and the CO 2.3\,$\mu$m bandheads in emission \citep{rei97a}, which is more commonly observed in EXors than FUors \citep[e.g.][]{lor06}. With its unique characteristics concerning outburst duration, repetition frequency, and spectroscopic diagnostics, V346\,Nor provides important new insights on the relation between the FUor and EXor phenomenon and might help to identify their triggering mechanism(s).

1607.03114

1607

1607.06394_arXiv.txt

{Very faint X-ray binaries appear to be transient in many cases with peak luminosities much fainter than that of usual soft X-ray transients, but their nature still remains elusive.} {We investigate the possibility that this transient behaviour is due to the same thermal/viscous instability which is responsible for outbursts of bright soft X-ray transients, occurring in ultracompact binaries for adequately low mass-transfer rates. More generally, we investigate the observational consequences of this instability when it occurs in ultracompact binaries.} {We use our code for modelling the thermal-viscous instability of the accretion disc, assumed here to be hydrogen poor. We also take into account the effects of disc X-ray irradiation, and consider the impact of the mass-transfer rate on the outburst brightness.} {We find that one can reproduce the observed properties of both the very faint and the brighter short transients (peak luminosity, duration, recurrence times), provided that the viscosity parameter in quiescence is slightly smaller (typically a factor of between two and fours) than in bright soft X-ray transients and normal dwarf nova outbursts, the viscosity in outburst being unchanged. This possibly reflects the impact of chemical composition on non-ideal MHD effects affecting magnetically driven turbulence in poorly ionized discs.} {}

Soft X-ray transients (SXTs) are a subclass of low mass X-ray binaries (LMXBs) which alternate between quiescent periods lasting years and bright outbursts which can reach the Eddington limit and last typically for weeks \citep[see e.g.][for reviews of the observations]{ts96,csl97,yy15}. These bright sources were discovered soon after the first X-ray satellites came into operation. With the advent of more sensitive X-ray telescopes such as \textit{Chandra}, \textit{XMM-Newton} and \textit{Swift}, fainter and fainter sources were discovered \citep[see e.g.][]{hww04,mpb05,pgb05,swd05} during observations of galactic fields and particularly the galactic centre. \citet{wzr06} then introduced two new classes of transient sources, the so-called faint X-ray transients, and very faint X-ray transients (VFXTs) that have peak luminosities in the range 10$^{36-37}$ erg s$^{-1}$, and 10$^{34-36}$ erg s$^{-1}$ respectively. These VFXTs are part of the very faint X-ray binaries group (VFXBs) which are mostly transients with quiescent luminosities below 10$^{33}$ erg s$^{-1}$ \citep{hbd15}. Despite the fact that they were discovered more than ten years ago, their nature still remains elusive, probably in part because of their faintness, and also because most of them are located in highly absorbed regions making the identification of optical counterparts difficult. It is in particular not clear that the groups of faint and very faint transients are homogeneous, and searching for a model accounting for properties of all VFXBs might be inappropriate. It is widely accepted that the transient nature of the classical soft X-ray transients (SXTs), which are much brighter in outburst than VFXBs and can reach luminosities of order of the Eddington limit, is due to the thermal/viscous instability of the accretion disc in regions where its effective temperature is of order of 6,000 -- 8,000 K \citep{kr98,dhl01}. This occurs when the mass-transfer rate from the secondary lies in some range $[\dot{M}_{\rm crit}^-(r_{\rm in}) - \dot{M}_{\rm crit}^+(r_{\rm out})]$ \citep{lasota-01}, where $r_{\rm in}$ and $r_{\rm out}$ are the inner and outer disc radius respectively, and $\dot{M}_{\rm crit}^-$ and $\dot{M}_{\rm crit}^+$ are the critical mass-transfer rates for the instability to occur at a given point in the disc ($\dot{M}_{\rm crit}^-$ is the maximum value for the disc to stay on the cold, stable branch, and $\dot{M}_{\rm crit}^+$ is the minimum value for the disc to stay on the hot branch). The observed luminosities of persistent and transient X-ray binaries show that this in indeed the case \citep{vp96,cfd12} if the accretion disc is X-ray self-irradiated. It is tempting to assume that faint and very faint outbursts could be due to the same instability that is responsible for the bright outbursts of SXTs. Such a model has to explain the faintness of the outbursts and take into account the low average mass accretion rate, which for VFXBs lies in the range $3 \times 10^{-13}$ -- $1 \times 10^{-10}$ M$_\odot$ yr$^{-1}$ \citep{dw09}. This low mass-transfer rate is directly related to the nature of the secondary and its evolutionary status, and, to a lesser extent, to the primary mass. \citet{k00} suggested that faint transients contain low mass secondaries that have evolved past the minimum orbital period. \citet{kw06} proposed later that VFXBs were either formed with a brown dwarf or a planetary companion, or could be the end products of massive primordial binaries which lead to systems containing an intermediate mass black hole, of the order of 1000 M$_\odot$. The latter explanation would of course not hold for VFXBs that exhibit type I X-ray bursts, and must therefore contain a neutron star. As noted by \citet{hbd15}, many of VFXBs have exhibited type I X-ray bursts, for example XMM J174457-2850.3 \citep{dwr14} and therefore do not harbor a black hole. \citet{mp13} proposed that the weak mass-transfer rates in VFXBs can be explained if these systems are detached, being in a similar state as cataclysmic variables in the period gap. This occurs when the secondary star becomes fully convective as a result of mass loss; angular momentum losses are reduced and the slightly bloated secondary shrinks, returning to thermal equilibrium; mass transfer can only occur via the faint secondary wind. The mass-transfer rate is difficult to estimate, but \citet{mp13} quote values of order of $10^{-16} - 10^{-14}$ M$_\odot$ yr$^{-1}$, which are much lower than the values of the order of $10^{-13}$ M$_\odot$yr$^{-1}$ inferred from observations of VFXTs \citep{dw10}. If the transient nature of the faint X-ray sources and VFXBs were to be attributed to the thermal/viscous disc instability, the observational difference between bright transients and VFXBs would have to be due either to a small accretion disc for the latter, (the critical mass-transfer rates $\dot{M}_{\rm crit}^+$ and $\dot{M}_{\rm crit^-}$ and the outburst peak luminosity scale as $r^{2.65}$), and/or to the fact that only a small fraction of the disc is involved in the instability. The smallest accretion discs are found in ultracompact binaries. These systems form a subclass of low-mass X-ray binaries (LMXBs) with very short orbital periods (typically less than $\sim$ 1hr). Unless the system has followed a very special evolutionary track, the secondary star is hydrogen poor in order to fit inside its Roche lobe, being either a degenerate helium star or the core of a carbon-oxygen or oxygen-neon-magnesium white dwarf \citep[see e.g.][]{N08}. Compared to hydrogen-rich binaries, this change in chemical composition significantly increases the temperature for which the disc becomes unstable, and hence the critical mass-transfer rates, but the global picture remains unchanged \citep{ldk08}. In this paper, we investigate in detail the nature and observational consequences of the disc instability model (DIM) applied to ultracompact binaries, in particular taking into account the effects of self-irradiation by accretion-produced X-rays. \cite{ldk08} estimated the stability conditions for hydrogen deficient ultracompact binaries, taking into account these effects, but have not produced light curves that can be directly compared with observations. This is needed in particular when one is interested in the faint end of the outburst luminosity distribution, for which the outermost parts of the disc may remain in the cold state during an outburst, and the peak luminosity is far from reaching the maximum luminosity one can derive using simple analytical estimates. \cite{kld12} did calculate the light curves predicted by the DIM in the case of degenerate helium secondaries, but, as they were interested in AM CVn stars, they did not consider disc irradiation. We show here that the observational properties (peak luminosities, duration, recurrence time) of many of the VFXB transients can be accounted for by the DIM provided, however, that the viscosity in quiescence is slightly smaller (typically a factor of between two and four) than in hydrogen--rich binaries. This does not imply that all VFXBs are ultracompact binaries -- some, such as for example \object{AX J1745.6-2901} which has an orbital period of 8.4 hr, are obviously not. Other systems, such as \object{1RXH J173523.7-354013} for which an H$\alpha$ line has been detected \citep[][note that this particular object is a persistent source, with an X-ray luminosity of order of $2 \times 10^{35}$ erg s$^{-1}$]{djt10} may contain a hydrogen rich secondary and are therefore not ultracompact. It does, however, mean that the short and faint outbursts observed in a number of VFXTs can be explained by the DIM when applied to ultracompact binaries. We leave for future work the much needed application of the DIM to the case where the donor star is a carbon-oxygen degenerate star, and in particular the confirmation that the C/O discs would behave in a manner similar to hydrogen rich discs.

We have shown that the disc instability model applied to ultracompact binaries produces outbursts whose characteristics are very similar to those of the faint and very faint outbursts observed in VFXBs, provided we use the same value of the viscosity parameter $\alpha_{\rm h}$ as in cataclysmic variables and in longer period X-ray binaries, and a smaller (factor a few) value of $\alpha_{\rm c}$ in quiescence. This is not surprising because the non-ideal-MHD effects that are likely affecting the magnetically driven turbulence in quiescent, poorly ionised discs \citep[e.g.][]{Bai14,Lesuretal14}, might depend on their chemical composition \citep[see, however,][]{kld12}. In any case, contrary to $\alpha_{\rm h}$, whose value is well determined from observations which consistently produced a value of $0.1 - 0.2$ \citep[see e.g.][]{KL12}, the value of $\alpha_{\rm c}$ does not have the same status because its observational determination is less direct than that of $\alpha_{\rm h}$. Various estimates give values from 0.0001 to 0.04 \citep[see e.g.][]{lasota-01}. The outburst faintness results from the small size of the accretion disc and from the low mass-transfer rate. In the case of a 1.4 M$_\odot$ neutron star accreting from a helium degenerate secondary, the secular mean of the mass-transfer rate is \citep{hnv12}: \begin{equation} \dot{M}_{\rm tr}= 6.66 \times 10^{13} P_{\rm hr}^{-5.32} \; \rm g s^{-1} . \end{equation} For the 0.7 hr orbital period considered here, this gives $\dot{M}_{\rm tr} = 4.4 \times 10^{14}$ gs$^{-1}$. Of course, the actual mass-transfer rate from the secondary need not be equal to the secular mean, but it is worth noting that the values used in this paper are within factors $\sim$ a few equal to the secular mean. We also predict that many of these systems with low mass-transfer rates could exhibit two classes of outbursts: those where the heat front is able to bring the whole disc in a hot state and are therefore relatively bright and last longer than the much fainter ones for which the heating front is unable to reach the outer disc edge. It is possible that these extremely faint outbursts are not detected as such because of their faintness -- in order to increase the sensitivity of short {\it Swift} exposures, data had to be added within time frames of 2-4 weeks --, and that they could be partly responsible for the luminosity variations observed during the intermediate state of very faint X-ray transients such as \object{XMM J174457-2850.3} \citep{dwr14}. One should, however, notice that the duty cycle we produce is small, and that it would be difficult to account for all of the observed luminosity variations observed during the intermediate state. Very faint outbursts can also be produced in long period binaries, provided that the mass-transfer rate is low enough. For the $\dot{M}_{\rm tr} = 10^{13}$ gs$^{-1}$ cases presented in Section 2.2.2, the heating front never reached distances larger than a third of the disc size, and the outer parts of the disc were therefore in a steady state on the cool branch; the outburst properties are therefore independent of the disc size, and the same light curves would have been obtained for the same mass-transfer rate for longer orbital periods. Very faint outbursts have been observed in several bright transients, such as for example \object{Aql X-1} \citep{ccd14}, \object{XTE J1701-462} \citep{f10}, \object{KS 1741-293} \citep{dw13} or SAX J1750.8-2900 \citep{wd13}. The possibility that these transients are of the same nature as those mentioned in this paper, with heat-front not able to propagate throughout the disc is interesting, but needs to be investigated in more details. This is clearly out of the scope of this paper. One should also notice that the truncated helium disc that we have considered here is stable when the mass-transfer rate is lower than $\dot{M}_{\rm crit}^-(r_{\rm in})$; for the parameters we have used here, this corresponds to X-ray luminosities of $10^{33} - 10^{34}$ erg s$^{-1}$. Such stable low luminosity systems have been observed in the case of AM CVn systems \citep[see e.g.][and references therein]{Bildstenetal06} and are predicted by the DIM \citep{smak83,kld12}. For a given orbital period, one therefore obtains the following sequence for decreasing the mass-transfer rates: for high $\dot{M}_{\rm tr}$, the system is stable on the hot branch; at lower $\dot{M}_{\rm tr} < \dot{M}_{\rm crit}^+(r_{\rm out})$, (relatively) bright outbursts are observed, during the entire disc is being brought in the hot state; for lower $\dot{M}_{\rm tr}$, one obtains a sequence of bright and faint outbursts during which the heating front does not reach the outer disc edge, and for still lower $\dot{M}_{\rm tr}$, only faint outbursts are obtained. Finally, for $\dot{M}_{\rm tr} < \dot{M}_{\rm crit}^-(r_{\rm in})$, the system is stable on the cold branch. In all cases, the outbursts are expected to be short, and depend mainly on $\alpha_{\rm c}$, and, for the outbursts able to reach the outer disc edge, on the orbital period. In all cases, faint outbursts are expected to be short if no additional ingredient, such as an increase of the mass-transfer rate due to the illumination of the secondary, is added. We have not taken into account the interaction of the neutron star magnetic field with the accreted matter. This magnetic field could be the reason for the disc truncation; it could also prevent accretion onto the neutron star via the propeller effect, thereby accounting for the observed transition between the millisecond pulsar state and the LMXB state \citep[see e.g.][]{pt15}. However, the recent observation of X-ray pulsations in the accreting millisecond X-ray pulsar \object{PSR J1023+0038} while the system was in quiescence with an X-ray luminosity of $3 \times 10^{33}$ erg s$^{-1}$ \citep{abp15} or of \object{XSS J12270-4859} which has similar properties \citep{pmb15} shows that even at those low mass--accretion rates, the propeller mechanism does not always prevent accretion. It is also worth noting that the light cylinder radius, equal to $4.8 \times 10^6 P_{\rm ms}$ cm, where $P_{\rm ms}$ is the spin period of the millisecond pulsar, is smaller than the truncation radius given by Eq. (\ref{eq:rtrunc}) for millisecond rotation periods. The effect of the neutron star magnetic field is thus more complex than usually assumed, and it is unlikely that simple recipes can be used to decide if the magnetic field is able to disrupt the accretion disc or not. Finally, one should keep in mind that the neither the VFXB nor the very faint X-ray transient categories are homogeneous. Apparently not all VFXBs are ultracompact and probably not all very faint outbursts can be explained by the DIM in ultracompact binaries.

1607.06394

1607

1607.01782_arXiv.txt

Standard approaches to Bayesian parameter inference in large scale structure assume a Gaussian functional form (chi-squared form) for the likelihood. This assumption, in detail, cannot % be correct. Likelihood free inferences such as Approximate Bayesian Computation (ABC) relax these restrictions and make inference possible without making any assumptions on the likelihood. Instead ABC relies on a forward generative model of the data and a metric for measuring the distance between the model and data. In this work, we demonstrate that ABC is feasible for LSS parameter inference by using it to constrain parameters of the halo occupation distribution (HOD) model for populating dark matter halos with galaxies. Using specific implementation of ABC supplemented with Population Monte Carlo importance sampling, a generative forward model using HOD, and a distance metric based on galaxy number density, two-point correlation function, and galaxy group multiplicity function, we constrain the HOD parameters of mock observation generated from selected ``true'' HOD parameters. The parameter constraints we obtain from ABC are consistent with the ``true'' HOD parameters, demonstrating that ABC can be reliably used for parameter inference in LSS. Furthermore, we compare our ABC constraints to constraints we obtain using a pseudo-likelihood function of Gaussian form with MCMC and find consistent HOD parameter constraints. Ultimately our results suggest that ABC can and should be applied in parameter inference for LSS analyses.

Cosmology was revolutionized in the 1990s with the introduction of likelihoods---% pro\-ba\-bil\-ities for the data given the theoretical model---% for combining data from different surveys and performing principled inferences of the cosmological parameters (\citealt{White:1996aa, Riess:1998aa}). Nowhere has this been more true than in cosmic microwave background (CMB) studies, where it is nearly possible to analytically evaluate a likelihood function that involves no (or minimal) approximations (\citealt{Oh:1999aa}, \citealt{Wandelt:2004aa}, \citealt{Eriksen:2004aa}, \citealt{planckI, planckII}). Fundamentally, the tractability of likelihood functions in cosmology flows from the fact that the initial conditions are exceedingly close to Gaussian in form (\citealt{planck_NG, planck_inflation}), and that many sources of measurement noise are also Gaussian (\citealt{Knox:1995aa, Leach:2008aa}). Likelihood functions are easier to write down and evaluate when things are closer to Gaussian, so at large scales and in the early universe. Hence likelihood analyses are ideally suitable for CMB data. In large-scale structure (LSS) with galaxies, quasars, and quasar absorption systems as tracers, formed through nonlinear gravitational evolution and biasing, the likelihood {\em cannot} be Gaussian. Even if the initial conditions are perfectly Gaussian, the growth of structure creates non-linearities which are non-Gaussian (see \citealt{Bernardeau:2002aa} for a comprehensive review). Galaxies form within the density field in some complex manner that is modeled only effectively (\citealt{Dressler:1980aa, Kaiser:1984aa, Santiago:1992aa, Steidel:1998aa}; see \citealt{somerville15} for a recent review). Even if the galaxies were a Poisson sampling of the density field, which they are not (\citealt{Mo:1996aa, Sommerville:2001aa, Casas-Miranda:2002aa}), it would be tremendously difficult to write down even an approximate likelihood function (\citealt{devpois}). The standard approach makes the strong assumption that the likelihood function for the data can be approximated by a pseudo-likelihood function that is a Gaussian probability density in the space of the two-point correlation function estimate. It is also typically limited to (density and) two-point correlation function (2PCF) measurements, assuming that these measurements constitute sufficient statistics for the cosmological parameters. As Hogg (in preparation) demonstrates, the assumption of a Gaussian pseudo-likelihood function cannot be correct (in detail) at any scale, since a correlation function, being related to the variance of a continuous field, must satisfy non-trivial positive-definiteness requirements. These requirements truncate function space such that the likelihood in that function space could never be Gaussian. The failure of this assumption becomes more relevant as the correlation function becomes better measured, so it is particularly critical on intermediate scales, where neither shot noise nor cosmic variance significantly influence the measurement. Fortunately, these assumptions are not required for cosmological inferences, because high-precision cosmological simulations can be used to directly calculate LSS observables. Therefore, we can simulate not just the one- or two-point statistics of the galaxies, but also any higher order statistics that might provide additional constraining power on a model. In principle, there is therefore no strict need to rely on these common but specious analysis assumptions as it is possible to calculate a likelihood function directly from simulation outputs. Of course, any naive approach to sufficiently simulating the data would be ruinously expensive. Fortunately, there are principled, (relatively) efficient methods for minimizing computation and delivering correct posterior inferences, using only a data simulator and some choices about statistics. In the present work, we use Approximate Bayesian Computation---ABC---which provides a \emph{rejection sampling} framework (\citealt{abcrejectionsampling}) that relaxes the assumptions of the traditional approach. ABC approximates the posterior probability distribution function (model given the data) by drawing proposals from the prior over the model parameters, simulating the data from the proposals using a forward generative model, and then rejecting the proposals that are beyond a certain threshold ``distance'' from the data, based on summary statistics of the data. In practice, ABC is used in conjunction with a more efficient sampling operation like Population Monte Carlo (PMC; \citealt{smc}). PMC initially rejects the proposals from the prior with a relatively large ``distance'' threshold. In subsequent steps, the threshold is updated adaptively, and samples from the proposals that have passed the previous iteration are subjected to the new, more stringent, threshold criterion (\citealt{abcpmc}). In principle, the distance metric can be any positive definite function that compares various summary statistics between the data and the simulation. In the context of astronomy, this approach has been used in a wide range of topics including image simulation calibration for wide field surveys (\citealt{abccosmology}), the study of the morphological properties of galaxies at high redshifts (\citealt{abcmorphology}), stellar initial mass function modeling (Cisewski et al. in preparation), and cosmological inference with with weak-lensing peak counts (\citealt{abcwl,abcwl2}), Type Ia Supernovae (\citealt{abcsn}), and galaxy cluster number counts (\citealt{cosmoabc}). In order to demonstrate that ABC can be tractably applied to parameter estimation in contemporary LSS analyses, we narrow our focus to inferring the parameters of a Halo Occupation Distribution (HOD) model. The foundation of HOD predictions is the halo model of LSS, that is, collapsed dark matter halos are biased tracers of the underlying cosmic density field (\citealt{press74, bond91, cooray_sheth2002}). The HOD specifies how the dark matter halos are populated with galaxies by modeling the probability that a given halo hosts $N$ galaxies subject to some observational selection criteria (\citealt{lemson99, seljak2000,scoccimarro2001,berlind_weinberg2002,zheng2005}). This statistical prescription for connecting galaxies to halos has been remarkably successful in reproducing the galaxy clustering, galaxy--galaxy lensing, and other observational statistics (\citealt{Rodriguez-Torres:2015aa, miyatake15}), and is a useful framework for constraining cosmological parameters (\citealt{vdb03, tinker05, cacciato13, more13}) as well as galaxy evolution models (\citealt{conroy09, Tinker:2011aa, leauthaud12, behroozi13, Tinker:2013aa}, Walsh et al. in preparation). More specifically, we limit our scope to a likelihood analysis of HOD model parameter space, keeping cosmology fixed. We forward model galaxy survey data by populating pre-built dark matter halo catalogs obtained from high resolution N-body simulations (\citealt{bolshoi,multidark}) using $\mathtt{Halotools}$\footnote{http://halotools.readthedocs.org} (\citealt{Hearin:2016aa}), an open-source package for modeling the galaxy-halo connection. Equipped with the forward model, we use summary statistics such as number density, two-point correlation function, galaxy group multiplicity function (GMF) to infer HOD parameters using ABC. In Section \ref{sec:method} we discuss the algorithm of the ABC-PMC prescription we use in our analyses. This includes the sampling method itself, the HOD forward model, and the computation of summary statistics. Then in Section \ref{sec:mock_obv}, we discuss the mock galaxy catalog, which we treat as observation. With the specific choices of ABC-PMC ingredients, which we describe in Section \ref{sec:abcpmc_spec}, in Section \ref{sec:abc_results} we present the results of our parameter inference using two sets of summary statistics, number density and 2PCF and number density and GMF. We also include in our results, analogous parameter constraints from the standard MCMC approach, which we compare to ABC results in detail, Section \ref{sec:abcvsmcmc}. Finally, we discuss and conclude in Section \ref{sec:discussion}.

\label{sec:discussion} Approximate Bayesian Computation, ABC, is a generative, simulation-based inference that can deliver correct parameter estimation with appropriate choices for its design. It has the advantage over the standard approach in that it does not require explicit knowledge of the likelihood function. It only relies on the ability to simulate the observed data, accounting for the uncertainties associated with observation and on specifying a metric for the distance between the observed data and simulation. When the specification of the likelihood function proves to be challenging or when the true underlying distribution of the observable is unknown, ABC provides a promising alternative for inference. The standard approach to large scale structure studies relies on the assumption that the likelihood function for the observables -- often two-point correlation function -- given the model has a Gaussian functional form. In other words, it assumes that the statistical summaries are Gaussian distributed. In principle to rigorously test such an assumption, a large number of realistic simulations would need to be generated in order to examine the actual distribution of the observables. This process, however, is prohibitively---both labor and computationally ---expensive. Therefore, our assumption of a Gaussian likelihood function remains largely unconfirmed and so unknown. Fortunately, the framework of ABC permits us to bypass any assumptions regarding the distribution of observables. Through ABC, we can provide constraints for our models without making the unexamined assumption of Gaussianity. With the ultimate goal of demonstrating that ABC is feasible for LSS studies, we use it to constrain parameters of the halo occupation distribution, which dictates the galaxy-halo connection. We begin by constructing a mock observation of galaxy distribution with a chosen set of ``true'' HOD model parameters. Then we attempt to constrain these parameters using ABC. More specifically, in this paper: \begin{itemize} \item We provide an explanation of the ABC algorithm and present how Population Monte Carlo can be utilized to efficiently reach convergence and estimate the posterior distributions of model parameters. We use this ABC-PMC algorithm with a generative forward model built with $\mathtt{Halotools}$, a software package for creating catalogs of galaxy positions based on models of the galaxy-halo connection such as the HOD. \item We choose $\ngalaxy$, $\xigg$ and $\gmf$ as observables and summary statistics of the galaxy position catalogs. And for our ABC-PMC algorithm, we specify a multi-component distance metric, uniform priors, a median threshold implementation, and an acceptance rate-based convergence criterion. \item From our specific ABC-PMC method, we obtain parameter constraints that are consistent with the ``true'' HOD parameters of our mock observations. Hence we demonstrate that ABC-PMC can be used for parameter inference in LSS studies. \item We compare our ABC-PMC parameter constraints to constraints using the standard Gaussian-likelihood MCMC analysis. The constraints we get from both methods are comparable in accuracy and precision. However, for our analysis using $\ngalaxy$ and $\gmf$ in particular, we obtain less biased posterior distributions when comparing to the ``true'' HOD parameters. \end{itemize} Based on our results, we conclude that ABC-PMC is able to consistently infer parameters in the context of LSS. We also find that the computation required for our ABC-PMC and standard Gaussian-likelihood analyses are comparable. Therefore, with the statistical advantages that ABC offers, we present ABC-PMC as an improved alternative for parameter inference.

1607.01782

1607

1607.06488_arXiv.txt

This study entails the third part of a global flare energetics project, in which {\sl Ramaty High-Energy Solar Spectroscopic Imager (RHESSI)} data of 191 M and X-class flare events from the first 3.5 yrs of the {\sl Solar Dynamics Observatory (SDO)} mission are analyzed. We fit a thermal and a nonthermal component to RHESSI spectra, yielding the temperature of the differential emission measure (DEM) tail, the nonthermal power law slope and flux, and the thermal/nonthermal cross-over energy $e_{\mathrm{co}}$. From these parameters we calculate the total nonthermal energy $E_{\mathrm{nt}}$ in electrons with two different methods: (i) using the observed cross-over energy $e_{\mathrm{co}}$ as low-energy cutoff, and (ii) using the low-energy cutoff $e_{\mathrm{wt}}$ predicted by the warm thick-target bremsstrahlung model of Kontar et al. {\bf Based on a mean temperature of $T_e=8.6$ MK in active regions we find low-energy cutoff energies of $e_{\mathrm{wt}} =6.2\pm 1.6$ keV for the warm-target model, which is significantly lower than the cross-over energies $e_{\mathrm{co}}=21 \pm 6$ keV. Comparing with the statistics of magnetically dissipated energies $E_{\mathrm{mag}}$ and thermal energies $E_{\mathrm{th}}$ from the two previous studies, we find the following mean (logarithmic) energy ratios with the warm-target model: $E_{\mathrm{nt}} = 0.41 \ E_{\mathrm{mag}}$, $E_{\mathrm{th}} = 0.08 \ E_{\mathrm{mag}}$, and $E_{\mathrm{th}} = 0.15 \ E_{\mathrm{nt}}$. The total dissipated magnetic energy exceeds the thermal energy in 95\% and the nonthermal energy in 71\% of the flare events, which confirms that magnetic reconnection processes are sufficient to explain flare energies. The nonthermal energy exceeds the thermal energy in 85\% of the events, which largely confirms the warm thick-target model.}

We undertake a systematic survey of the global energetics of solar flares and coronal mass ejections (CME) observed during the SDO era, which includes all M and X-class flares during the first 3.5 years of the SDO mission, covering some 400 flare events. This project embodies the most comprehensive survey about various forms of energies that can be detected during flares, such as the dissipated magnetic energy, the thermal energy, the nonthermal energy, the radiative and conductive energy, and the kinetic energy of associated CMEs. Two studies have been completed previously, containing statistics on magnetic energies (Aschwanden Xu, and Jing 2014; Paper I), and thermal energies (Aschwanden et al.~2015; Paper II). In this study we focus on the third part of this ``global flare energetics project", which entails the statistics of nonthermal energies in hard X ray-producing electrons that are observed in hard X-rays and gamma-rays, using data from the Ramaty High-Energy Solar Spectroscopic Imager (RHESSI) spacecraft (Lin et al.~2002). The quantitative measurement of nonthermal energies in solar flares allows us some tests of fundamental nature. One concept or working hypothesis is that all primary energy input in solar flares is provided by dissipation of free magnetic energy, for instance by a magnetic reconnection process, which supplies energy for secondary processes, such as for acceleration of charged particles and heating of flare plasma. The accelerated (nonthermal) particles either escape from the flare site into interplanetary space, or more likely precipitate down to the chromosphere where they subsequently become thermalized and radiate in hard X-rays and gamma rays, according to the thick-target bremsstrahlung model (Brown 1971). In this picture we expect that the total nonthermal energy $E_{\mathrm{nt}}$ (in electrons and ions) produced in flares should not exceed the dissipated magnetic (free) energy $E_{\mathrm{mag}}$, but on the other hand should yield an upper limit on the thermal energy $E_{\mathrm{th}}$ inferred from the soft-X-ray and EUV-emitting plasma. Alternative mechanisms to the thick-target model envision thermal conduction fronts (e.g., Brown et al.~1979) or direct heating processes (e.g., Duijveman et al.~1981). In the previous two papers we proved the inequality $E_{\mathrm{mag}} > E_{\mathrm{th}}$, for which we found an energy conversion ratio of $E_{\mathrm{th}}/E_{\mathrm{mag}} \approx 0.02-0.40$ (Paper II), which is about an order of magnitude higher than estimated in a previous statistical study (Emslie et al.~2012), where an {\sl ad hoc} value (30\%) of the ratio of the free magnetic energy to the potential field energy was estimated. In this Paper III we investigate the expected inequalities $E_{\mathrm{mag}} > E_{\mathrm{nt}} > E_{\mathrm{th}}$. If these two inequalities are not fulfilled, it could be attributed to insufficient accuracy of the energy measurements, or alternatively may question the correctness of the associated low-energy cutoff model, the applied magnetic reconnection models, or the efficiency of the electron thick-target bremsstrahlung model. Such an outcome would have important consequences in our understanding of solar flare models and the related predictability of the most extreme space weather events. The measurement of nonthermal energies in solar flares requires a spectral fit of the hard X-ray spectrum in the energy range of $\varepsilon \approx 10-30$ keV (Aschwanden 2007), from spectral data as they are available from the HXRBS/SMM, BATSE/CGRO, or RHESSI instrument. Since the total nonthermal energy contained in a flare requires integrations over the temporal and spectral range, the largest uncertainty of this quantity comes from the assumed low-energy cutoff, because it cannot be directly measured due to the strong thermal component that often dominates the spectrum at $\varepsilon \lapprox 20$ keV during solar flares (for a review see Holman et al.~2011). In a few cases, low-energy cutoffs of the nonthermal spectrum could be determined by regularized inversion methods at $e_c=20-40$ keV (Kasparova et al.~2005), $e_c \approx 20$ keV (Kontar and Brown 2006), and $e_c=13-19$ keV (Kontar, Dickson, and Kasparova 2008). For the 2002 July 23 flare, Holman et al.~(2003) deduced upper limits to low-energy cutoffs by determining the highest values consistent with acceptable spectral fits. Sui et al.~(2007) deduced the low-energy cutoff in a flare from the combination of spectral fits and the time evolution of the X-ray emission in multiple energy bands. Sui et al.~(2007) deduced low-energy cutoffs for several flares with relatively weak thermal components (``early impulsive flares'') from spectral fits, with values ranging from $15-50$ keV. In the late peak of a multi-peaked flare, Warmuth et al.~(2009) inferred low-energy cutoff values exceeding 100 keV, but this unusually high value could possibly be explained also by high-energy electrons that accumulate by trapping after the flare peak (Aschwanden et al.~1997). Using a novel method of {\bf differentiating} nonthermal electrons by their time-of-flight delay from thermal electrons by their thermal conduction time delay, a thermal-nonthermal crossover energy of $e_c=18.0 \pm 3.4$ keV (or a range of $e_c = 10-28$ keV) was established for the majority (68\%) of 65 analyzed flare events (Aschwanden 2007). Statistical measurements of nonthermal flare energies have been calculated from HXRBS/SMM data (Crosby et al.~1993), or from RHESSI data (Hannah et al.~2008; Christe et al.~2008; Emslie et al.~2012). The low-energy cutoff was taken into account by assuming a fixed energy cutoff of $e_c=25$ keV (Crosby et al.~1993), a fixed spectral slope of $\gamma=-1.5$ below the thermal-nonthermal cross-over energy $e_{\mathrm{co}}$ (Hannah et al.~2008), or by adopting the largest energy $e_c$ that still produces a goodness-of-fit with $\chi^2 \approx 1$ for the nonthermal power law fit (Emslie et al.~2012). Low-energy cutoffs for microflares were estimated in the range of $e_c \approx 9-16$ keV, with a median of 12 keV (Hannah et al.~2008), using a numerical integration code of Holman (2003). The statistical study of Emslie et al.~(2012) provides a comparison between nonthermal energies $E_{\mathrm{nt}}$, thermal energies $E_{\mathrm{th}}$, and dissipated magnetic energies $E_{\mathrm{mag}}$, yielding mean (logarithmic) ratios of $E_{\mathrm{th}} \approx 0.005\ E_{\mathrm{mag}}$ and $E_{\mathrm{nt}} \approx 0.03\ E_{\mathrm{mag}}$. These results conform to the expected inequalities, but the magnetic energies $E_{\mathrm{mag}}$ were actually not measured in the study of Emslie et al.~(2012), and most likely were overestimated by an order of magnitude (Paper I). The dissipated magnetic energies $E_{\mathrm{mag}}$ were for the first time quantitatively measured in Paper I, by automated tracing of coronal flare loops from AIA/SDO images and by forward-fitting of a nonlinear force-free magnetic field (NLFFF) model based on the vertical current approximation (Aschwanden 2013, 2016). The content of this paper consists of a theoretical model to estimate the low-energy cutoff and the nonthermal energy (Section 2), a description of the data analysis method (Section 3), the results of the data analysis of 191 M and X-class flare events observed with RHESSI (Section 4), a discussion of the results (Section 5), and conclusions (Section 6).

The energy partition study of Emslie et al.~(2012) was restricted to 38 large solar eruptive events (SEE). In a more comprehensive study on the global flare energetics we choose a dataset that contains the 400 largest (GOES M and X-class) flare events observed during the first 3.5 years of the SDO era. Previously we determined the dissipated magnetic energies $E_{\mathrm{mag}}$ in these flares based on fitting the {\sl vertical-current approximation of a nonlinear force-free field (NLFFF)} solution to the loop geometries detected in EUV images from SDO/AIA, a new method that could be applied to 177 events with a heliographic longitude of $\le 45^\circ$ (Paper I). We also determined the thermal energy $E_{\mathrm{th}}$ in the soft X-ray and EUV-emitting plasma during the flare peak times based on a multi-temperature differential emission measure DEM forward-fitting method to SDO/AIA image pixels with spatial synthesis, which was applicable to 391 events (Paper II). In the present study we determined the nonthermal energy $E_{\mathrm{nt}}$ contained in accelerated electrons based on spectral fits to RHESSI data using the OSPEX software, which was applicable to 191 events. The major conclusions of the new results emerging from this study are: \begin{enumerate} \item{The (logarithmic) mean energy ratio of the nonthermal energy to the total magnetically dissipated flare energy is found to be $E_{\mathrm{nt}}/E_{\mathrm{mag}}=0.41$, with a logarithmic standard deviation corresponding to a factor of $\approx 8$, which yields an uncertainty $\sigma/\sqrt{N}=0.41/\sqrt{191}=0.03$ for the mean, i.e., $E_{\mathrm{nt}}/E_{\mathrm{mag}}=0.41\pm0.03$. The majority ($\approx 85\%$) of the flare events fulfill the inequality $E_{\mathrm{nt}}/E_{\mathrm{mag}} < 1$, which suggests that magnetic energy dissipation (most likely by a magnetic reconnection process) provides sufficient energy to accelerate the nonthermal electrons detected by bremsstrahlung in hard X-rays. Our results yield an order of magnitude higher electron acceleration efficiency than previous estimates, i.e., $E_{\mathrm{nt}}/E_{\mathrm{mag}} =0.03\pm 0.005$ (with $N=37$, Emslie et al.~2012).} \item{The (logarithmic) mean of the thermal energy $E_{\mathrm{th}}$ to the nonthermal energy $E_{\mathrm{nt}}$ is found to be {\bf $E_{\mathrm{th}}/E_{\mathrm{nt}}=0.15$, with a logarithmic standard deviation corresponding to a factor of $\approx 7$. The fraction of flares with a thermal energy being smaller than the nonthermal energy, as expected in the thick-target model, is found to be the case for $\approx 85\%$ only. Therefore, the thick-target model is sufficient to explain the full amount of thermal energy in most flares, in the framework of the warm-target model. The cross-over method shows the opposite tendency, but we suspect that the cross-over method over-estimates the low energy cutoff and under-estimates the nonthermal energies. Previous estimates yielded a similar ratio, i.e., $E_{\mathrm{th}}/E_{\mathrm{nt}}=0.15$ (Emslie et al.~2012).}} \item{A corollary of the two previous conclusions is that the thermal to magnetic energy ratio is $E_{\mathrm{th}}/E_{\mathrm{mag}}=0.08$. A total of $95\%$ flares fulfils the inequality $E_{\mathrm{nt}}/E_{\mathrm{mag}} < 1$, indicating that all thermal energy in flares is supplied by magnetic energy. Previous estimates were a factor of 17 lower, i.e., $E_{\mathrm{th}}/E_{\mathrm{mag}}=0.0045$ (Emslie et al.~2012), which would imply a very inefficient magnetic to thermal energy conversion process.} \item{The largest uncertainty in the calculation of nonthermal energies, the low-energy cutoff, is found to yield different values for two used methods, i.e., $e_{\mathrm{wt}}=6.2 \pm 1.6$ keV for the warm thick-target model, versus $e_{\mathrm{co}}=21 \pm 6$ keV for the thermal/nonthermal cross-over method. The calculation of the nonthermal energies is highly sensitive to the value of the low-energy cutoff, which strongly depends on the assumed (warm-target) temperature.} \item{The flare temperature can be characterized with three different definitions, for which we found the following ($67\%$-standard deviation) ranges: $T_{\mathrm{AIA}} \approx 3-14$ MK for the AIA DEM peak temperature, $T_w \approx 20-30$ MK for the emission measure-weighted temperatures, and $T_{R} \approx 17-36$ MK for the RHESSI high-temperature DEM tails. The median ratios are found to be $T_{\mathrm{AIA}}/T_w=0.31$ and $T_{R}/T_w=0.90$. {\bf The mean active region temperature evaluated from DEMs with AIA, $T_e=8.6$ MK, is used to estimate the low-energy cutoff $e_c$ of the nonthermal component according to the warm-target model, i.e., $e_c \approx \delta (k_B T_{R})$. The low-energy cutoff $e_c$ of the nonthermal spectrum has a strong functional dependence on the temperature $T_{R}$.}} \end{enumerate} In summary, our measurements appear to confirm that the magnetically dissipated energy is sufficient to explain thermal and nonthermal energies in solar flares, which strongly supports the view that magnetic reconnection processes are the primary energy source of flares. The nonthermal energy, which represents the primary energy source of the thick-target model, {\bf is sufficient to explain the full amount of thermal energies in 71\% of the flares, according to the novel warm-target model (Kontar et al.~2011). However, the derived nonthermal energies are highly dependent on the the assumed temperature in the warm-target plasma, for which a sound physical model should be developed (see for instance Appendix A and B), before it becomes a useful tool to estimate the low-energy cutoff of nonthermal energy spectra.} Future studies of this global flare energetics project may also quantify additional forms of energies, such as the kinetic energy in CMEs, and radiated energies in soft X-rays, EUV, and white-light (bolometric luminosity). \bigskip

1607.06488

1607

1607.05442_arXiv.txt

The halo of the Milky-Way circumgalactic gas extends up to the virial radius of the Galaxy, $\sim 250$~kpc. The halo properties may be deduced from X-ray spectroscopic observations and from studies of the ram-pressure stripping of satellite dwarf galaxies. The former method is more precise but its results depend crucially on the assumed metallicity of the circumgalactic gas; the latter one does not need these assumptions. Here, the information from both approaches is combined to constrain observationally the gas metallicity and density as functions of the galactocentric distance. It is demonstrated that the \blue two kinds of data could be reconciled if the \black metallicity decrease\blue{}d \black \red to $Z\sim 0.1Z_{\odot}$ \black in the outer parts of the extended halo. \red The corresponding gas density profile is rather flat, falling as $r^{-(0.45 \dots 0.75)}$ at large galactocentric distances $r$. \black

\label{sec:intro} Recently, considerable attention has been attracted to studies of gas coronae of galaxies, that is of reservoirs of gas extending up to the galaxies' virial radii. This circumgalactic gas represents, thanks to the large volume it fills, a substantial contribution to the mass budget of a galaxy. This gaseous corona, or extended halo, of the Milky Way has attracted particular interest because of the ``missing-baryon'' problem, see e.g.\ \citet{DM}, the apparent lack of baryons in our Galaxy compared to the amount expected, on average, from cosmology. On the other hand, interactions of cosmic rays with this circumgalactic gas have been considered as a possible source of important contributions to the diffuse gamma-ray \citep{Hooper} and neutrino \citep{Aha} backgrounds. In the Milky Way, this reservoir of gas reveals itself in observations in two ways. First, the ram pressure of the gas strips dwarf satellite galaxies, whose orbits lay within the corona, from their own gas \citep{dwarfs250kpc}. Second, the presence of the hot gas may be seen in X-ray spectra, either as zero-redshift absorption lines for extragalactic sources, or as emission lines in the blank-sky spectrum, see e.g.\ \citet{gas1,MB2013,1412.3116}. Taken at face value, the gas density profiles derived by these two methods are inconsistent with each other. However, the spectroscopic approach is based on observations of spectral lines of oxygen, which is only a tracer of the full amount of gas. As a result, the gas density obtained by the spectroscopic method is very sensitive to unknown chemical composition of the gas, usually encoded in its metallicity $Z$. The aim of this work is to depart from simplified ad hoc assumptions about the metallicity of the Galactic corona and to use spectroscopic and ram-pressure results jointly, which allows us to constrain values and profiles of density and metallicity of the circumgalactic gas simultaneously, so that the agreement between all data is maintained. The rest of the paper is organized as follows. In Sec.~\ref{sec:obs}, observational constraints on the density of circumgalactic gas are discussed in detail. In particular, in Sec.~\ref{sec:X-ray}, X-ray spectroscopic results are discussed and their dependence on the assumptions about metallicity is recalled. Sec.~\ref{sec:ram} discusses constraints from ram-pressure stripping of the Milky-Way satellites; a combined fit of the most precise of these bounds is presented. Other constraints are briefly mentioned in Sec.~\ref{sec:DM}. Section~\ref{sec:comb} contains the main results of the paper and presents a combination of the constraints, allowing to determine both the density and the metallicity of circumgalactic gas in a joint fit by means of statistical marginalization. These results are discussed and compared to previous works in Sec.~\ref{sec:concl}.

\label{sec:concl} The results of this work eliminate the apparent discrepancy, see Fig.~\ref{fig:2contours}, in parameters of the Milky-Way circumgalactic gas estimated from X-ray spectroscopy and from studies of ram-pressure stripping of Galactic dwarf satellites. A combined analysis of the observational constraints, performed here, determines the range of the allowed metallicity profiles of the gas residing up to 250~kpc from the Galactic Center (Figs.~\ref{fig:ABmarg}, \ref{fig:met-profiles}). Not surprisingly, the metallicity decreases considerably in the outer parts of the halo \red with respect to the inner part. The profile in the outer part is, however, fairly flat. \black Importantly, we obtained constraints on the parameters of the gas density distribution from a combination of all available data (Figs.~\ref{fig:density-marg}, \ref{fig:density-profiles}). Compared to the profile of \citet{1412.3116}, the best-fit one is flatter, resulting in higher gas density in the peripheral parts of the Galactic corona and, consequently, in larger total gas mass. This agrees well with recent simulations \citep{HalfHidden} and ultraviolet O{\small VI} observations \citep{1602.00689} indicating that the X-ray observations may underestimate the total amount of circumgalactic gas by a factor of two. It is interesting to compare the total gas mass with the ``missing baryon'' mass of the Galaxy, since the halo of circumgalactic gas was suggested as an explanation of the mismatch between the Milky-Way and cosmological average baryon content \citep{DM}. One can see from Fig.~\ref{fig:density-marg}, where a line corresponding to the required ``missing-baryon'' mass is shown, that our results support this explanation. \red It is also interesting to see how our density profile agrees with estimates of the temperature and the luminosity of the Galactic X-ray halo. To this end, we note the relation between the slope parameter $ \beta$ in Eq.~(\ref{Eq:n_e}), the velocity dispersion of galactic objects $\sigma$ and the gas temperature $T$, \begin{equation} \beta=\frac{\mu m_p \sigma^2}{kT}, \label{Eq:T} \end{equation} see e.g.\ \citet{NFW}, where $\mu$ is the mean atomic mass per particle, $m_{p}$ is the proton mass and $k$ is the Boltzmann constant. The velocity dispersion is consistent with $\sigma \sim 100$~km/s in the inner $\sim 80$~kpc \citep{1304.5127}, but several observations indicate significant decrease of $\sigma$ at large $r$ \citep{a-p/0506102,0910.2242}. We use $\sigma=90$~km/s in the following estimates. The estimated values of $T$ are shown in the right scale of Fig.~\ref{fig:T}. \begin{figure} \begin{center} \includegraphics[width=0.95 \columnwidth]{fig-T.jpg} \caption{ \label{fig:T} \red Estimates of the halo temperature $T$ and X-ray luminosity $L_{X}$ versus the density profile parameters. The red diamond and the red contour represent, respectively, the best-fit point and the 68\% CL contour obtained in this work. The right scale represents $T$ estimated from $\beta$. Blue horizonthal lines give the $T$ median value (full line, $2.22 \times 10^{6}$~K) and interquantile range (dashed lines, $0.63 \times 10^{6}$~K) from \citet{1306.2312}. Thin gray lines bound the range $L_{X}=(2\dots 3)\times 10^{39}$~erg/s favoured by \citet{ApJ_485_125,a-p/9710144}. See the text for details and important notes. \black } \end{center} \end{figure} Given the approximate nature of these estimates, our density profiles are in a good agreement with observational constraints on the X-ray temperature of the halo gas, for instance, those by \citet{1306.2312}, shown in Fig.~\ref{fig:T}. We also calculate the total X-ray luminosity of the halo, $L_{x}$, by making use of Eqns.~(17)-- (19) of \citet{1412.3116}. The metallicity enters there through the cooling function \citep{ApJ_88_253} which we take for [Fe/H]$=-1$, corresponding to the most part of the halo, cf.\ Fig.~\ref{fig:met-profiles}. Observations point to $L_{X}\sim (2\dots 3)\times 10^{39}$~erg/s \citep{ApJ_485_125,a-p/9710144}, the range shown in Fig.~\ref{fig:T} as well, again in a good agreement with our preferred parameters for the density profile. However, these \blue temperature-related \black estimates should be considered with caution, because of several reasons. Firstly, the constraints discussed in this paper are relevant for the outer part of the halo, where observational information on X rays is scarce. Secondly, the estimates assume that the halo is isothermal while some studies point to the opposite \citep{0906.1532}. Thirdly, they assumed constant metallicity and velocity dispersion. Finally, the quantitative values of temperature and luminosity depend on the values of poorly known parameters to which the results of the present paper are insensitive, like $r_{c}$ and $\sigma$\blue, for which we have very little data to work with\black. Therefore, these considerations and results presented in Fig.~\ref{fig:T} should be considered only as a demonstration of the qualitative agreement of our model with observational data on $T$ and $L_{X}$. \blue Indeed, e.g., the replacement of the velocity-dispersion temperature in Eq.~(\ref{Eq:T}) by the rotation-curve based temperature estimate would change the gas temperature by a factor of $\sim 2$, indicating a factor of $\sim 2$ higher $\beta$; however, the accuracy of Eq.~(\ref{Eq:T}) is of the same order. Clearly, the beta models themselves might be very crude tools for modelling of the possibly structured circumgalactic gas medium, but relaxing the constant-metallicity assumption and removal of the discrepancies we discuss here are necessary first steps towards understanding of this interesting part of the Galaxy. The gas density profile obtained in this work is considerably flatter than the total mass density profile of the halo, in qualitative agreement with simulations by \black \citet{Hooper}. At galactocentric distances $\lesssim 40$~kpc, the profile of \citet{Hooper} deviates from the beta profile, Eq.~(\ref{Eq:n_e}), towards higher densities. Unfortunately, this range of distances is not controlled by our approach; therefore, higher densities are not experimentally excluded there.

1607.05442

1607

1607.02054_arXiv.txt

We investigate the properties of $z$=2.23 \ha\ and \oiii$\lambda$5007 emitters using the narrow-band-selected samples obtained from the High-$z$ Emission Line Survey (HiZELS: \citealt{sobral13}). We construct two samples of the \ha\ and \oiii\ emitters and compare their integrated physical properties. We find that the distribution of stellar masses, dust extinction, star formation rates (SFRs), and specific SFRs, is not statistically different between the two samples. When we separate the full galaxy sample into three subsamples according to the detections of the \ha\ and/or \oiii\ emission lines, most of the sources detected with both \ha\ and \oiii\ show ${\rm log(sSFR_{UV})}$$\gtrsim$-9.5. The comparison of the three subsamples suggests that sources with strong \oiii\ line emission tend to have the highest star-forming activity out all galaxies that we study. We argue that the \oiii\ emission line can be used as a tracer of star-forming galaxies at high redshift, and that it is especially useful to investigate star-forming galaxies at $z$$>$3, for which \ha\ emission is no longer observable from the ground.

\label{intro} Emission lines from regions ionized by hot, young massive stars are useful as indicators of star formation in distant galaxies. Imaging observations with narrow-band (NB) filters, which can capture redshifted strong emission lines, are a powerful method to construct a star-forming galaxy sample at a particular redshift (e.g. \citealt{bunker95,malkan96,moorwood00,geach08,sobral09,sobral13,tadaki13,an14,stroe15}). The \ha\ emission line is one of the best tracers of star formation because it is less affected to dust extinction than the ultraviolet (UV) light and the relation between star-formation rates (SFRs) and \ha\ luminosities has been well calibrated in the local Universe (e.g. \citealt{hopkins03}). This seems to hold at higher redshift. In addition, \ha\ selection has the advantage of recovering the full population of star-forming galaxies \citep{oteo15}. However, the redshift range for \ha\ selection is limited to $z$$<$2.8 because \ha\ emission is no longer easily observed beyond $z$$\sim$2.8 with ground-based telescopes. Space telescopes, such as the {\it Spitzer}, have been used for the \ha\ emission line galaxy survey at higher redshift, $z$$\gtrsim$4.0, but only using broad band photometry (e.g. \citealt{shim11,smit15}), which is therefore only sensitive to the highest equivalent width lines. In order to investigate star-forming galaxies at $z$$>$2.8 with NB imaging observations, it is necessary to use other emission lines at shorter wavelengths than \ha, such as \oii$\lambda$3727, \hb, and \oiii$\lambda5007$. With \oii, \hb, and \oiii\ emission lines, we can reach up to $z$$\sim$5.2, 3.7 and 3.6, respectively, from the ground \citep{khostovan15}. \oii\ and \hb\ are relatively weak lines and so it is more difficult to observe them at higher redshift. There is also evidence of higher \oiii/\oii\ line ratios at high redshifts, and a potential decline in \oii\ equivalent width \citep{khostovan16}. On the other hand, strong \oiii\ detections that are comparable to \ha\ from high-redshift star-forming galaxies have been reported by recent near-infrared (NIR) spectroscopic observations (e.g. \citealt{holden14,masters14,steidel14,shimakawa15,shapley15}). Such strong \oiii\ emission indicates extreme interstellar medium (ISM) conditions in high-redshift galaxies, and this is likely to be due to their lower metallicities and/or higher ionization parameters (e.g. \citealt{nakajima14}). Also, \oiii\ emission in the rest-frame optical is less sensitive to dust extinction than the UV light. In these respects, it is expected that the \oiii\ emission line can be used to select star-forming galaxies at $z$$\sim$3--3.6, corresponding to $\sim$1--1.5 billion years before the peak epoch of galaxy formation and evolution at $z$$\sim$2 (e.g. \citealt{hopkins06,khostovan15}). Some studies have been constructing \oiii\ (+ \hb) emitter samples at $z$$>$3, and have investigated their star-forming activity or the evolution of the luminosity function \citep{teplitz99,maschietto08,labbe13,smit14,suzuki15, khostovan15,khostovan16}. Star-forming activity and other physical properties of galaxies at $z$$>$3 have also been investigated with UV-selected galaxies, such as Lyman Break Galaxies (e.g. \citealt{stark09,gonzalez10,reddy12,stark13,tasca14}). By using \oiii\ emission as a tracer at $z$$>$3, it is expected that we can obtain further understanding about galaxy formation and evolution before the peak epoch of galaxy assembly. However, there are possible biases resulting from the use of \oiii\ emission as a star-forming indicator. For example, the \oiii\ emission line originates from ionized regions not only caused by hot, young massive stars in star-forming regions but also by active galactic nuclei (AGNs; e.g. \citealt{zakamska04}). Due to the shorter wavelength of \oiii\ (5007\AA) with respect to \ha\ (6563\AA), samples of \oiii-selected galaxies may be inherently biased against dusty systems in comparison to an \ha-selected sample. As already mentioned above, galaxies with strong \oiii\ emission might be biased towards galaxies with lower metallicities and/or higher ionization states, resulting in a potential bias towards galaxies with lower stellar masses given the well-known mass--metallicity relation of star-forming galaxies (e.g. \citealt{tremonti04,erb06,stott14,troncoso14}). At $z$$\sim$0 and $z$$\sim$1.5, the selection biases between \ha\ and \oiii\ have been investigated using the SDSS galaxies and FMOS-COSMOS galaxies \citep{silverman14} by Juneau et al. (in preparation, private communication). Their samples are selected based on \ha\ and \oiii\ luminosities. \citet{mehta15} investigated the relation between \ha\ and \oiii\ luminosity for galaxies at $z$$\sim$1.5 in the {\it HST}/WFC3 Infrared Spectroscopic Parallel Survey (WISP; \citealt{atek10}), and derived the \ha--\oiii\ bivariate luminosity function at $z$$\sim$1.5. \citet{sobral12} and \citet{hayashi13} investigated the relation between the NB-selected \ha\ and \oii\ emission line galaxies at $z$$\sim$1.47. They discussed the lack of redshift evolution of the \oii/\ha\ ratio from $z$$\sim$0 to 1.5 \citep{sobral12}, and also found that the \oii-selected galaxies tend to be biased towards less dusty galaxies with respect to the \ha-selected galaxies \citep{hayashi13}. At $z$$>$2, some studies have already performed the comparisons between the samples of star-forming galaxies selected by different optical emission lines, such as \ha\ and Ly$\alpha$, or by other selection methods (e.g. \citealt{oteo15,hagen15,matthee16,shimakawa16}). However, the comparison of the physical quantities between \ha\ and \oiii-selected galaxy samples has not been done yet at $z$$>$2. Such a comparison is necessary in order to accurately interpret results from \oiii\ surveys at $z$$>$3. In this study, we use the NB-selected \oiii\ and \ha\ emission line galaxies at $z$=2.23, obtained by HiZELS (the High-$z$ Emission Line Survey; \citealt{best13,sobral09,sobral12,sobral13,sobral14}), a large NB imaging survey. The \oiii\ and \ha\ emission lines are observed using the \nbh\ and \nbk\ filters, respectively. This combination of NB filters allows for the creation of a suitable sample to investigate possible selection biases between the \oiii\ and \ha\ emission line galaxies at high redshift. We compare the integrated physical quantities, such as stellar masses, dust extinction, and star formation rates (SFRs), and investigate whether there are any systematic differences between these physical quantities. Some galaxies are detected with both the \oiii\ and \ha\ emission lines. We investigate the physical properties of the galaxies depending on the detectability of their \oiii\ and \ha\ emission lines. This paper is organized as follows: In \S \ref{data}, we briefly introduce the NB imaging survey, HiZELS, and describe how the \oiii\ and \ha\ emission line galaxies at $z$$\sim$2 are selected. We also present the method for deriving the integrated physical quantities. Then, we show our results in \S \ref{results}. We present the relationship between stellar masses and SFRs for the two emitter samples, and compare the number distribution of the global physical quantities between the \ha\ and \oiii\ emitters. Moreover, we divide our full sample into three subsamples according to the detections of the \ha\ and/or \oiii\ emission lines, and compare the distributions of physical quantities among the three subsamples. We summarize this study in \S \ref{summary}. We assume the cosmological parameters of $\Omega_{\rm m}=0.3$, $\Omega_{\Lambda}=0.7$, and $H_{\rm 0}=70 \ [{\rm km\ s^{-1}Mpc^{-1}}]$. Throughout this paper all the magnitudes are given in AB magnitude system \citep{oke83}, and the Salpeter initial mass function (IMF; \citealt{salpeter55}) is adopted for the estimation of the stellar masses and SFRs \footnote{Stellar mass estimated assuming the Salpeter IMF can be scaled to those assuming the Chabrier \citep{chabrier03} and Kroupa \citep{kroupa02} IMF by dividing by a factor of $\sim$ 1.7 and 1.6, respectively \citep{pozzetti07,marchesini09}.}.

\label{summary} We use the NB-selected galaxy catalog at $z$=2.23 obtained by the HiZELS project, and construct the two galaxy samples of \ha\ and \oiii\ emission line galaxies by applying the same line flux limit. We derive the global physical properties of these emitters, and compare the number distribution of a stellar mass, dust extinction ($A_{\rm FUV}$), ${\rm SFR_{UV}}$, and ${\rm sSFR_{UV}}$ between the two samples. The resulting $p$-values from a KS-test indicates that the \ha\ and \oiii\ emitters are drawn from the same parent population. The two galaxy populations cover almost the same ranges of the integrated properties at $z$$\sim$2. We also divide the whole sample into three subsamples, namely, the galaxies detected with either \ha\ or \oiii\ alone, and the galaxies detected with both lines. Again, a KS-test does not show any significant differences among the three subsamples, except for the dual emitters, which tend to be biased to higher ${\rm sSFR_{UV}}$ as compared to the other two subsamples. It is indicated that the strong \oiii\ emission lines are likely to be related to high star formation activities (and thus high ionization parameters) of star-forming galaxies at $z$$\sim$2. Note, however, the \oiii\ and \ha\ emitters used in this study could harbor low luminosity AGNs, especially the \oiii-single-emitters with low sSFR as discussed in \S \ref{exception}, and spectroscopic observations are necessary to confirm the presence of AGNs. In summary, the \oiii\ emitters trace almost the same galaxy populations as the \ha\ emitters at $z$$\sim$2, and therefore we argue that the \oiii\ emission line can be used as an indicator of normal star-forming galaxies at high redshifts. Our results support the importance and the effectiveness of \oiii\ emitter surveys at $z$$\gtrsim$3, where \ha\ emission is no longer effectively observed from the ground.

1607.02054

1607

1607.07447_arXiv.txt

We present chemical abundances derived from high-resolution Magellan/MIKE spectra of the nine brightest known red giant members of the ultra-faint dwarf galaxy Reticulum~II. These stars span the full metallicity range of Ret~II ($-3.5 < \mbox{[Fe/H]} < -2$). Seven of the nine stars have extremely high levels of $r$-process material ([Eu/Fe]$\sim 1.7$), in contrast to the extremely low neutron-capture element abundances found in every other ultra-faint dwarf galaxy studied to date. The other two stars are the most metal-poor stars in the system ([Fe/H] $< -3$), and they have neutron-capture element abundance limits similar to those in other ultra-faint dwarf galaxies. We confirm that the relative abundances of Sr, Y, and Zr in these stars are similar to those found in $r$-process halo stars but $\sim 0.5$\,dex lower than the solar $r$-process pattern. If the universal $r$-process pattern extends to those elements, the stars in Ret~II display the least contaminated known $r$-process pattern. The abundances of lighter elements up to the iron peak are otherwise similar to abundances of stars in the halo and in other ultra-faint dwarf galaxies. However, the scatter in abundance ratios is large enough to suggest that inhomogeneous metal mixing is required to explain the chemical evolution of this galaxy. The presence of low amounts of neutron-capture elements in other ultra-faint dwarf galaxies may imply the existence of additional $r$-process sites besides the source of $r$-process elements in Ret~II. Galaxies like Ret~II may be the original birth sites of $r$-process enhanced stars now found in the halo.

\label{s:intro} Ultra-faint dwarf galaxies (UFDs) probe extreme astrophysical regimes. They are the faintest and most metal-poor galaxies known \citep{Kirby08,Kirby13b}. Their high velocity dispersions imply they are the most dark matter dominated galaxies \citep{Simon07,Strigari08,Simon11}, making them attractive targets for indirect dark matter searches (e.g., \citealt{DrWag15}). The bulk of their star formation occurs before reionization \citep{Brown14}, and they may be important sources of ionizing photons \citep{Weisz14b,Wise14}. The initial mass function in UFDs differs from more massive galaxies \citep{Geha13}. Most importantly for our current purpose, UFDs provide a coherent environment in which to probe the earliest stages of nucleosynthesis and chemical evolution \citep{Frebel12,Karlsson13,Ji15}. Reticulum II (henceforth {\RetII}) is a UFD recently discovered in the Dark Energy Survey \citep{Koposov15a,Bechtol15}. Its velocity dispersion and metallicity spread confirm it to be a galaxy, and it is one of the most metal-poor galaxies known \citep{Simon15,Walker15,Koposov15b}. At only {$\sim$}30 kpc away, it contains stars within the reach of high-resolution spectroscopy for abundance analysis. Until recently, nearly all UFD stars observed with high-resolution spectroscopy displayed unusually low neutron-capture element abundances compared to halo star abundances (\AB{X}{Fe}$\lesssim -1$) (e.g., \citealt{Frebel10b,Frebel14,Koch13}). However, \citet{Ji16b} and \citet{Roederer16b} reported that seven of the nine stars they observed in {\RetII} have highly enhanced neutron-capture abundances (\AB{Eu}{Fe}$ \sim 1.7$). Moreover, the relative abundances of the elements heavier than barium match the scaled solar {\rproc} pattern \citep{Sneden08}, confirming that the universality of this nucleosynthesis process holds for stars in the faintest dwarf galaxies (also see \citealt{Aoki07b}). Metal-poor stars with this level of {\rproc} enhancement (\AB{Eu}{Fe} $> 1$, or {\rII} stars, \citealt{Christlieb04}) are only rarely found in the halo \citep{Barklem05,Roederer14d}. The striking 2-3 orders of magnitude difference between the neutron-capture element content of {\RetII} and that of the other UFDs is clear evidence that a single rare and prolific {\rproc} event is responsible for nearly all neutron-capture material in {\RetII} \citep{Ji16b}. In addition to usual questions about the formation history of UFDs and possible signatures of the first stars, this galaxy provides a tremendous opportunity to study the origin of the {\rproc} elements. \citet{Roederer16b} presented the first high resolution abundance measurements of elements lighter than barium in four {\RetII} stars. They found the abundances of Sr, Y, and Zr in the three {\rproc}-rich {\RetII} stars were similar to those of the {\rII} star {\CSSneden}. They also found that the abundances of the sub-iron-peak elements were generally consistent with halo star abundances at similar metallicities, implying that the source of {\rproc} elements in {\RetII} either produced none of these elements or produced them in similar amounts as core-collapse supernovae. \citet{Roederer16b} also found abundance variations for different stars with similar {\feh}, which suggests that metals are not uniformly mixed into the galaxy's gas reservoir. Accounting for this inhomogeneous metal mixing is important for using chemical abundances to understand the formation of this galaxy (e.g., \citealt{Webster16}). Here, we report the complete chemical abundance patterns for the nine Reticulum II stars considered by \citet{Ji16b}, including the four investigated by \citet{Roederer16b}. Our stars span the entire metallicity range of {\RetII} \citep{Simon15}. In Section~\ref{s:methods} we describe the observations and abundance analysis. The abundance patterns are reported in Section~\ref{s:abunds}. In Section~\ref{s:nuclear} we discuss implications for nuclear astrophysics and the {\rproc} site. In Section~\ref{s:stargalform} we consider possibilities for using this galaxy to understand early star and galaxy formation. We conclude in Section~\ref{s:concl}.

1607.07447

1607

1607.02781_arXiv.txt

Since their discovery, the Quintuplet proper members (QPMs) have been somewhat mysterious in nature. Originally dubbed the ``cocoon stars" due to their cool featureless spectra, high-resolution near-infrared imaging observations have shown that at least two of the objects exhibit ``pinwheel'' nebulae consistent with binary systems with a carbon-rich Wolf-Rayet star and O/B companion. In this paper, we present 19.7, 25.2, 31.5, and 37.1 $\mu$m observations of the QPMs (with an angular resolution of 3.2-3.8") taken with the Faint Object Infrared Camera for the SOFIA Telescope (FORCAST) in conjunction with high-resolution ($\sim$0.1-0.2") images at 8.8 and 11.7 $\mu$m from the Thermal-Region Camera Spectrograph (TReCS). DUSTY models of the thermal dust emission of two of the four detected QPMs, Q2 and Q3, are fitted by radial density profiles which are consistent with constant mass loss rates ($\rho_d \propto r^{-2}$). For the two remaining sources, Q1 and Q9, extended structures ($\sim 1"$) are detected around these objects in high-resolution imaging data. Based on the fitted dust masses, Q9 has an unusually large dust reservoir ($\mathrm{M_d}=1.3^{+0.8}_{-0.4}\times 10^{-3} \mathrm{M_{\odot}}$) compared to typical dusty Wolf-Rayet stars which suggests that it may have recently undergone an episode of enhanced mass loss.

Five very luminous infrared sources reside in the young, massive Quintuplet cluster in the Galactic center (Okuda et al. 1990). These five objects, for which the cluster is named, are referred to as the Quintuplet proper members (QPMs). The QPMs are sources of great interest, given their enigmatic nature exhibiting high infrared luminosities ($\sim10^{4}-10^5$ $\mathrm{L}_\odot$), cool spectral energy distributions ($400-800$ K), and previous claims of featureless near-infrared (IR) spectra (Moneti et al. 2001; Okuda et al. 1990; Nagata et al. 1990). Figer et al (1999) first proposed that the QPMs might be extremely dusty carbon-rich Wolf-Rayet stars (WCLd), but were unable to confirm this hypothesis without the identification of spectral features in the near-IR. From observations of the QPMs with the Infrared Space Observatory Camera (ISOCAM), Moneti et al. (2001) further explored this possibility, noting that the high IR luminosities ($\sim10^{4}-10^5$ $\mathrm{L}_\odot$) and small geometric dust covering fraction ($\sim$0.1) would be consistent with typical WC star luminosities; however, the lack of near-IR features in the J-band was still somewhat of a mystery. More recent multi-epoch, high-resolution near-IR images of two of the QPMs taken by the Keck 1 telescope have revealed ``pinwheel'' shaped plumes that are typically associated with colliding-wind binary systems (Tuthill et al. 2006). In the interpretation of Tuthill et al. (2006), the QPMs are late-type carbon-rich Wolf Rayet stars with an O/B star binary companion that form dust in the wind-wind collision front that is swept outward by the Wolf-Rayet winds (Tuthill et al. 1999). The classification of the QPMs as dust-enshrouded, evolved massive stars is consistent with the age of the Quintuplet cluster $\sim$3.3-3.6 Myrs (Liermann et al. 2012), which would be difficult to reconcile if the QPMs were young stellar objects (YSOs) as was originally suggested (Okuda et al. 1990). Although Tuthill el al. (2006) provide convincing evidence of the `pinwheel' nature of two of the QPMs: Q2 and Q3 (GCS 3-2 and GCS 4), the identification of the remaining members (Q1, Q4, and Q9 or alternatively, GCS 3-4, GCS 3-1, and GCS 3-3) was less certain. However, recent near-IR spectra taken by Geballe et al. (2014) clearly show broad emission line features associated with WC stars from Q1 and Q4 but not Q9. The two near-IR pinwheel stars, Q2 and Q3, appear to share somewhat similar properties. Blackbody fits to the SEDs of each object from Moneti et al. (2001) yield temperatures for Q2 and Q3 of 650 and 625 K. From the near-IR images, the systems appear to have, at least roughly, a similar geometry and are close to face-on (Tuthill et al. 2006). Q2 shows non-thermal radio emission and potentially weak x-ray emission (Lang et al. 2005; Law \& Yusef-Zadeh 2004) which is a telltale sign of a colliding wind system (Monnier et al. 2002). Q3 has no detected radio emission but exhibits stronger x-ray emission than Q2 (Law \& Yusef-Zadeh 2004). The remaining QPMs show no detectable radio or x-ray emission. Previous high-spatial-resolution images ($\sim 1''$) at 8.7 and 11.7 $\mu$m taken with the Palomar telescope show that the nebulae associated with Q9 and Q1 appear larger in physical extent than the other QPMs at low intensity levels and also display characteristically cooler spectra (400 and 500 K) compared with Q3 and Q2 (Moneti et al. 2001). Conversely, Q4 appears warmer (725 K) and more compact than Q2 and Q3. In this paper, we present 19.7, 25.2, 31.5, and 37.1 $\mu$m observations of the QPMs taken by FORCAST aboard the Stratospheric Observatory for Infrared Astronomy (SOFIA). These observations provide important information on the mid-IR portion of the spectral energy distribution (SED) of the QPMs which can be modeled to provide estimates of dust production and mass loss rates of these objects. Additionally, we improve on previous high-resolution mid-IR imaging with 8.8 and 11.7 $\mu$m images from TReCS in full-pupil and Sparse Aperture Masking (SAM) mode to study the morphology of the dust reservoirs in these objects. In a broader context, we aim to understand the effects of binarity in these systems as a study of massive evolved stars through their mass loss history. In a recent paper by Sana et al. (2012), it is suggested that $\sim$70\% of massive stars will exchange mass or merge with a binary companion over the course of their lifetime (see also Sana et al 2014). Thus, characterization of massive evolved binary systems is vital to understand the end states of massive stars. The objective of this analysis is to better classify these sources in terms of their thermal dust emission and offer possible explanations for their various similarities and differences.

In this paper, infrared observations taken with SOFIA/FORCAST and Gemini/TReCS were used to study the thermal dust emission from the QPMs. DUSTY models of the SEDs were used to determine properties of the dust reservoirs present in each object. Best-fit model parameters indicate that the near-IR pinwheel stars (Q2 and Q3) follow a $r^{-2}$ density profile, indicating a constant mass loss rate, while Q1 and Q9 show departures from a $r^{-2}$ density profile. High-resolution imaging shows large ($\sim$1'') extended structures associated with the latter two objects. Detailed modeling coupled with SAM observations favor the presence of larger dust grains than expected based on theoretical calculations in two of three objects where we have measurements. This is consistent with previous studies of dust grain sizes in other dusty WC stars. Based on the observed dust mass in these systems and the available radio mass loss rates, we derive a higher dust fraction than the ISM value in Q2 which is consistent with similar measurements of WR112. The massive dust reservoir in Q9 (M=$1.3\times10^{-3} \mathrm{M}_{\odot}$) is difficult to reconcile with the mass loss rates and wind speeds of typical dusty WC stars. The large dust reservoir and unusual density profile would suggest some kind of recent enhanced mass loss with $\dot{\mathrm{M}}\sim10^{-3}\mathrm{M}_{\odot}/$yr. Study of the literature on the WCLd population shows that periods of such enhanced mass loss in these kinds of objects are not typical, which possibly points to the highly unstable nature of massive stars near the end of their lives and the short timescales of the observable dust components in these types of systems. Further study of Q9 is needed to understand the nature of this very interesting object. \emph

1607.02781

1607

1607.03567.txt

We constrain the warm dark matter (WDM) particle mass with the observations of cosmic reionization and CMB optical depth. We suggest that the GWs from stellar mass black holes (BHs) could give a further constraint on WDM particle mass for future observations. The star formation rates (SFRs) of Population I/II (Pop I/II) and Population III (Pop III) stars are also derived. If the metallicity of the universe have been enriched beyond the critical value of $Z_{\rm crit}=10^{-3.5}Z_{\odot}$, the star formation shift from Pop III to Pop I/II stars. Our results show that the SFRs are quite dependent on the WDM particle mass, especially at high redshifts. Combing with the reionization history and CMB optical depth derived from the recent \emph{Planck} mission, we find that the current data requires the WDM particle mass in a narrow range of $1\kev \lesssim m_{\rm x}\lesssim 3\kev$. Furthermore, we suggest that the stochastic gravitational wave background (SGWB) produced by stellar BHs could give a further constraint on the WDM particle mass for future observations. For $m_{\rm x}=3 \kev$ with Salpeter (Chabrier) initial mass function (IMF), the SGWB from Pop I/II BHs has a peak amplitude of $\Omega_{\rm GW}\approx2.8\times 10^{-9}~(5.0\times 10^{-9})$ at $f= 316 {\rm Hz}$, while the GW radiation at $f<10$Hz is seriously suppressed. For $m_{\rm x}=1 \kev$, the SGWB peak amplitude is the same as that of $m_{\rm x}=1\kev$, but a little lower at low frequencies. Therefore, it is hard to constrain the WDM particle mass by the SGWB from Pop I/II BHs. To assess the detectability of GW signal, we also calculate the signal to noise ratio (SNR), which are $\rm SNR=37.7~ (66.5)$ and $27~(47.7)$ for $m_{\rm x}=3\kev$ and $m_{\rm x}=1\kev$ for Einstein Telescope (ET) with Salpeter (Chabrier) IMF, respectively. The SGWB from Pop III BHs is seriously dependent on the WDM particle mass, the GW strength could be an order of magnitude different and the frequency band could be two times different for $m_{\rm x}=1\kev$ and $m_{\rm x}=3\kev$. Moreover, the SGWB from Pop III BHs with $m_{\rm x}=1\kev$ could be detected by LISA for one year of observation, but can not for $m_{\rm x}=3\kev$.

The astrophysical and cosmological probes have confirmed that baryons constitute only some $16\%$ of the total matter in the Universe. The rest of the mass is in the form of `dark matter' (DM). The nature of DM particles is poorly understood, as they do not interact with baryons. Many indirect searches have been carried out, including searching for $\gamma$-ray signals at the Galactic center, in nearby galaxies, and the diffuse $\gamma$-ray background \citep{Ackermann12, Ackermann14, The Fermi LAT collaboration15}. However, none of them could provide robust evidence for the observation of DM. The GeV $\gamma$-ray excess from the Galactic center could be a signal of DM annihilation, but still can not be confirmed \citep{Daylan15, Zhou15}. Among various DM candidates, the most popular candidate is the weakly interacting massive particles (WIMPs; like the neutralino), which have mass in GeV range \citep{Jungman96, Bertone05, Hooper07, Feng10}. The WIMPs are non-relativistic at the epoch of decoupling from the interacting particles and have negligible free-streaming velocities. Therefore, they are `cold', called cold dark matter (CDM). In CDM scenario, `halos' formed in small clumps, and then merged together into larger and massive objects. Galaxies formed in these halos are because of the cooling of atomic hydrogen \citep[H;][]{Tegmark97} or molecular hydrogen \citep[${\rm H_2}$;][]{Ciardi00, Haiman00}. On large cosmological scales (from the range $\sim1\Gpc$ down to $\sim10\mpc$), CDM paradigm has great success in explaining the observed universe and reproducing the luminous structures \citep{Fixsen96, Borgani01, Lange01, Cole05, Tegmark06, Benson10, Wang13, Hinshaw13, Slosar13, Planck Collaboration14, Wei16}. However, on small scales ($\lesssim1\mpc$), there are still some discrepancies between the CDM paradigm and observations: (a) the core-cusp problem \citep{Navarro97, Subramanian00}. CDM simulations predict a cusp-core DM halo, whereas the observations find them cored \citep{Salucci12}; (b) too big to fail problem \citep{Boylan-Kolchin12}. CDM simulations predict a central DM density significantly higher than the observation that allowed; and (c) the `missing satellite problem'. N-body simulations based on the CDM paradigm predict a number of subhalos larger than that of satellites found in our Galaxy \citep{Klypin99, Moore99, Papastergis11}. Many methods have been proposed to solve these small scale problems, such as modifying the nature of DM from the CDM paradigm \citep{Hu00, Spergel00, Su11, Menci12}, adding supernova feedback effect in simulation \citep{Weinberg02, Mashchenko06, Governato10, Pontzen14}, and considering the interplay between DM and baryons during the formation of the galaxy \citep{El-Zant01, Tonini06, Pontzen14}. However, these methods are insufficient to solve all the above problems. Alternatively, a more possible solution to these small scale problems is the warm dark matter (WDM) scenario, with DM particle mass in $\kev$ range. The candidates are sterile neutrinos \citep{Dodelson94, Abazajian01, Abazajian06, Shaposhnikov06, Boyarsky09, Kusenko09, Abazajian12} and gravitinos \citep{Kawasaki97, Gorbunov08}. WDM particles are lighter than CDM particles, so they could remain relativistic for longer time in the early universe and retain a non-negligible velocity dispersion. They are more easy to free-stream out from small scale perturbations, and suppress the formation of subhalos \citep{Bode01, Lovell14}. The most powerful test for WDM scenario is the high-redshift universe. A number of works have been done to constrain the WDM particle mass ($m_{\rm x}$). For example, \cite{Kang13} gave a lower limit of $m_{\rm x}\gtrsim 0.75\kev$ by reproducing the stellar mass functions and Tully-Fisher relation for $0<z<3.5$ galaxies. \cite{Viel13} used Lyman-$\alpha$ flux power spectrum measured from high-resolution spectra of 25 quasars to obtain a lower limit of $m_{\rm x}\gtrsim 3.3\kev$. \cite{de Souza13} used high-redshift ($z>4$) gamma-ray bursts to constrain $m_{\rm x}\gtrsim 1.6-1.8 \kev$. \cite{Dayal15b} constrained $m_{\rm x}\gtrsim 2.5 \kev$ by comparing the semi-analytic merger tree based framework for high-redshift ($z\simeq 5-20$) galaxy formation with reionization indicators. \cite{Lapi15} gave a narrow constraint of $2<m_{\rm x}<3 \kev$ by combining the measurements of the galaxy luminosity functions out $z\sim 10$ from \emph{Hubble Space Telescope} (HST) with the reionization history of the universe from the \emph{Planck} mission. \cite{Pacucci13} constrained $m_{\rm x}\gtrsim 1 \kev$ by using the number density of $z\approx 10$ lensed galaxies. Given that structures are formed hierarchically and WDM scenario smears out the power on small scale, the number density of the smallest halos (or galaxies) at high redshift will be strongly decreased, and then the SFR. Especially the SFRs of Pop III stars and high-redshift Pop I/II stars, because they are firstly formed in these small halos \citep{Barkana01}. Pop III stars are the massive stars with masses $\gtrsim 100 M_\odot$ \citep[e.g.,][]{Bromm99, Abel00, Nakamura01}, which are formed in metal-free gas. The deaths of Pop III stars lead to the metal enrichment of intergalactic medium (IGM) via supernova feedback, and subsequently the formation of Pop I/II stars \citep[the cricitical metallicity is $10^{-3.5}Z_{\rm \odot}$;][]{Ostriker96, Madau01, Bromm03, Furlanetto03}. The mass of Pop I/II stars is in the range of $0.1\sim 100 M_{\odot}$. The first light from Pop III stars brought the end of the cosmic dark ages, and then the universe began to reionize. Recent observation from \emph{Planck} mission measured the integrated CMB optical depth with $\tau=0.066^{+0.013}_{-0.013}$ \citep[with the constraint from \emph{Planck} TT+low polarization+lesing+BAO;][]{Planck Collaboration15}, and most of the observations show that the universe was fully reionized at redshift $z\simeq 6$ \citep{Chornock13, Treu13, Pentericci14, Schenker14, McGreer15}. These measurements gauge the level of the reionization history from the high-redshift stars. Furthermore, as the high-redshift stars formed in small halos are greatly affected by the halo number density, their formation rates could provide a indirect test on the WDM scenario. On September 14, 2015 the Advanced LIGO observed the gravitational-wave event GW150914 \citep{Abbott16a}. The observed signal is consistent with a black-hole binary waveform with the component masses of $m_1=36^{+5}_{-4}\msun$ and $m_2=29^{+4}_{-4}\msun$, which demonstrates the existence of stellar-mass black holes massive than $25 \msun$. The second GW candidate GW 151226 was observed by the twin detectors of the Advanced LIGO on December 26, 2015 \citep{Abbott16b}. The inferred initial BH masses are $14.2_{-3.7}^{+8.3}\msun$ and $7.5_{-2.3}^{+2.3}\msun$, and the final BH mass is $20.8_{-1.7}^{+6.1}\msun$. The decay of the waveform at the final period are also observed, which are consistent with the damped oscillations of a black hole relaxing to a stationary Kerr configuration. The collapses of Pop I/II or Pop III stars into black holes (BHs) could also release gravitational waves \citep[GWs;~][]{Buonanno05, Sandick06, Suwa07, Pereira10, Ott13, Yang15}, which is dominated by `quasi-normal ringing' of a perturbed black hole. Therefore, it is expected that the Advanced LIGO could also observe this kind of gravitational wave radiations. In this paper, we will calculate the SGWB from BH `ringing', which relates with the de-excitation of the BH quasi-normal modes. Because the SGWB is quite dependent on the SFR, it could be used to constrain the WDM particle mass indirectly. Several GW detectors are operating or planed in future: advanced VIRGO and LIGO working at $\approx\rm 10 Hz-3kHz$, the Einstein Telescope (ET) with the sensitive frequency of $1-100 \rm Hz$, the Laser Interferometer Space Antenna \footnote{http://lisa.nasa.gov/} (LISA) covering the frequency range of $10^{-4}-0.1 \rm Hz$, the Decihertz Interferometer Gravitational wave Observatory\footnote{http://universe.nasa.gov/program/vision.html} (DECIGO)\citep{Kudoh06}, and the Big Bang Observer (BBO) operating in the range $0.01-10$ Hz. Therefore, GW signal from BHs ringing will open a new window for the restriction of the WDM particle mass. This paper is organized as follows. In section 2, we describe the hierarchical formation scenario in the framework of WDM paradigm. In section 3, we construct the SFRs of Pop I/II and Pop III stars, and compare them with the recent observations. In section 4, we constrain the WDM particle mass with the CMB optical depth and the reionization history. In section 5, we calculate the SGWBs from Pop I/II and Pop III BHs. Finally, conclusion and discussion are given in section 6. Throughout this paper, we adopt the standard flat cosmology with cosmological parameters $\Omega_\Lambda=0.72, \Omega_{\rm m}=0.28, \Omega_{\rm b}=0.046, H_0=70\km\s^{-1}\mpc^{-1}$, and $\sigma_8=0.82$.

Although the CDM paradigm has great success in explaining the large scale structure of the universe, it still have some problems on small scales. An alternatively WDM paradigm could ease these problems by employing the $\kev$ WDM particles. In this paper, we calculate the SFRs of Pop I/II and Pop III stars in the framework of WDM paradigm. By using a self-consistent method, we reproduce the SFR at low redshift. We find that the high-redshift SFR is sensitively dependent on the WDM particle mass, especially for Pop III stars. By comparing the model-predicted SFR with the observation, we constrain the WDM particle mass with $m_{\rm x}>1\kev$, shown as the black lines in Figure \ref{Figure1}. We also calculated the metal enrichment history of IGM, and the transition from Pop III to Pop I/II stars is consistent with the previous results \citep[e.g.,][]{Yang15}, e.g., from $z\sim 10-17$ for $m_{\rm x} \sim 1-3\kev$. By considering that the metallicity of IGM does not exceed $Z_{\sun}$, the WDM particle mass should less than $3\kev$. Combing with the CMB optical depth from \emph{Planck} with $\tau=0.066^{+0.013}_{-0.013}$ and the ionization fraction $Q_{\rm H_{\rm II}}$ from recent observations, we found that the the WDM particle mass should in the range of $1\kev \lesssim m_{\rm x}\lesssim3\kev$, where we have assumed a constant escape fraction of ionizing photons \citep[e.g.,][]{Schultz14, Dayal15a}. However, many works suggest that the escape fraction should be redshift dependent \citep[e.g.,][]{Siana10, Blanc11, Hayes11, Kuhlen12, Dijkstra14}. By considering an evolving escape fraction, we found a more tight constraint $1\kev< m_{\rm x}<2\kev$. Finally, recent observation of GW150914 and GW151226 inspires a great interest in the field of GW. Therefore, we calculated the SGWBs form Pop I/II and Pop III BHs. Our results show that the SGWB from Pop I/II BHs is not sensitive to the WDM particle mass, and it could be detected by the ET telescope. However, it is impossible to constrain the WDM particle mass by the SGWB from Pop I/II BHs, because they show little difference for different $m_{\rm x}$. For Pop III stars, the SGWB is quite dependent on the WMD particle mass. The peak SGWB amplitude with $m_{x}=1\kev$ is an order of magnitude higher than $m_{x}=3\kev$. The corresponding SNR are $1.76~(1.7)$ and 0.33 (0.35) for SIMF (CIMF), respectively, which is distinguishable for LISA. Moreover, the SGWBs are derived by assuming a maximum GW generation efficiency of $\epsilon_{\rm GW}=7\times 10^{-4}$. Combing with the ET observation of SGWB from Pop I/II BHs, we could give a constraint on $\epsilon_{\rm GW}$. Therefore, a further constraint of $m_{\rm x}$ (or Pop III SFR) could be given by the observation of LISA. On the other hand, the SGWB from Pop III BHs is also quite dependent on the star formation efficiency of $f_2$. Therefore, a lower efficiency of $f_2$ will make it hard to constrain the WDM particle mass by the observation of LISA. Anyway, the large difference of SGWBs from Pop III BHs for different $m_{\rm x}$ will make it possible to constrain the WDM particle mass in future.

1607.03567

1607

1607.05732_arXiv.txt

There is still much debate surrounding how the most massive, central galaxies in the local universe have assembled their stellar mass, especially the relative roles of in-situ growth versus later accretion via mergers. In this paper, we set firmer constraints on the evolutionary pathways of the most massive central galaxies by making use of empirical estimates on their abundances and stellar ages. The most recent abundance matching and direct measurements strongly favour that a substantial fraction of massive galaxies with \ms \textgreater $3 \times 10^{11}\ M_\odot$ reside at the centre of clusters with mass \mh \textgreater $3 \times 10^{13}\ M_\odot$. Spectral analysis supports ages \textgreater 10 Gyrs, corresponding to a formation redshift \zf \textgreater $2$. We combine these two pieces of observationally-based evidence with the mass accretion history of their host dark matter haloes. We find that in these massive haloes, the stellar mass locked up in the central galaxy is comparable to, if not greater than, the total baryonic mass at \zf. These findings indicate that either only a relatively minor fraction of their present-day stellar mass was formed in-situ at \zf, or that these massive, central galaxies form in the extreme scenario where almost all of the baryons in the progenitor halo are converted into stars. Interestingly, the latter scenario would not allow for any substantial size growth since the galaxy's formation epoch either via mergers or expansion. We show our results hold irrespective of systematic uncertainties in stellar mass, abundances, galaxy merger rates, stellar initial mass function, star formation rate and dark matter accretion histories.

The mechanisms both to form and evolve massive early-type galaxies are still highly debated. Early-type galaxies (ETG) dominate the red population of objects in the observed bimodal distribution of galaxy colours \citep{Baldry2004,Cassata2008}. However, how ETGs form, evolve and transition in the colour-mass plane remains unclear \citep[][and references therein]{Woo2015}. Significant effort has been put into probing the evolution of the most massive galaxies both observationally and theoretically. From the observational side, a renewed interest in this field emerged with the advent of large and deep galaxy surveys such as SDSS \citep{Ahn2014}, COSMOS \citep{Scoville2007}, BOSS \citep{BOSS2013} and CANDELS \citep{Grogin2011}. A key result of these surveys is the observed evolution of the size-mass relation of ETGs. Massive, early type galaxies are observed to be progressively more compact at higher redshift as compared to galaxies with the same mass in the local universe \citep{VanDokkum2010,Huertas-Company2013}. This size evolution could be triggered by later mergers \citep[e.g.,][]{Naab2009,Shankar2013}, gas accretion \citep{Dekel2009} and/or nearly adiabatic expansion due to quasar mode feedback and/or stellar winds \citep{Fan2008,Damjanov2009}. However, despite significant progress, issues such as ``progenitor bias'' \citep[e.g.,][]{VanDokkum1996, Saglia2010,Carollo2013,Shankar2015}, the role of environment \citep{Poggianti2006,Shankar2013,Delaye2014,Shankar2014a,Stringer2015} and observational systematics such as cosmic variance and stellar mass estimates \citep{Marchesini2009,Bernardi2013}, hinder a clear interpretation of the observed mass and size evolution, especially at the extreme high-mass end of the stellar mass function \cite[e.g.,][]{Marchesini2014, Shankar2014a,Leauthaud2016,Bernardi2016}. From the theoretical side, evolutionary models have not converged on one clear picture of massive galaxy evolution. Semi-analytic models are an effective way of probing the diverse physical processes that are believed to drive galaxy formation and evolution \citep{Cole2000,Baugh2006,Guo2011,Benson2012,Lacey2015}. These models, sometimes based on very different input assumptions, can offer degenerate solutions in reproducing a handful of key statistical properties such as the galaxy stellar mass function \citep[see review discussion in][]{Mo2010}. Broadly speaking, semi-analytic models predict two conflicting evolutionary pathways, one where dry mergers dominate the evolution \citep{DeLucia2011,Gonzalez2011,Guo2011,Shankar2013,Wilman2013} and one where in-situ processes are more important \citep{Lapi2011,Ragone-Figueroa2011,Chiosi2012,Merlin2012,Posti2014}. Hierarchical merger models predict that massive galaxies have assembled most of their final stellar mass via a sequence of mergers following their host dark matter haloes \cite[e.g.,][]{Naab2009,Shankar2009,VanDokkum2010,Guo2011,Shankar2013,Montes2014}. Indeed massive galaxies must have merged at some point as tidal tails and concentric shells are observed around massive, local galaxies \citep{Duc2015}. The rate at which galaxies merge is usually observationally inferred by looking for pairs of galaxies in close spatial proximity \cite[e.g.,][]{Hopkins2010}. However, this rate is non-trivial to quantify \citep[e.g.,][]{Conselice2014} as a number of systematics may effect the result, from the assumptions on dynamical friction timescales, to the determination of spectroscopic pairs. Consequently, the true role of mergers in shaping massive galaxies remains still uncertain. In-situ galaxy evolutionary models claim instead that massive galaxies formed and assembled most of their final stellar mass in strong bursts of star formation at high redshifts. These starbursts can have star formation rates as high as several thousands of solar masses per year \citep{Chapman2005}. After the starburst has quenched, possibly induced by an efficient quasar mode feedback, the galaxy is assumed to evolve almost passively until the present day \citep{Granato2004,Granato2006,Carollo2013,Zolotov2015}. ETGs are observed to be enhanced in alpha-elements relative to their iron content which is evidence for short bursts of intense star formation \citep{Thomas2005,Pipino2009,Conroy2014,Citro2016}. Hydrodynamical zoom simulations \citep[e.g.,][]{Hirschmann2012} have converged on the idea that there are two phases to massive galaxy evolution where in-situ star formation dominates the early assembly and mergers become more important at lower redshifts \citep{Naab2009,Oser2010}. Hydrodynamical simulations in a full cosmological box continue to support this two-stage evolutionary patten at least for the most massive galaxies \citep{Hirschmann2012,Torrey2015,Welker2015}. However, the relative roles of in-situ versus late assembly remains poorly constrained observationally. In recent years, a number of semi-empirical approaches have been put forward to more securely probe and constrain the possible evolutionary pathways of massive galaxies. For example, \cite{VanDokkum2010}, \cite{Marchesini2014} and \cite{Huertas-Company2015} have adopted number conservation techniques to track the putative main progenitors of massive galaxies. Other techniques are based on continuity equation models for the stellar population \citep[e.g.,][]{Peng2010,Aversa2015}. Also, \cite{Lidman2012} and \cite{Shankar2015} followed the main progenitor track of the host haloes to identify potential proto-galaxies as progenitors. All of these semi-empirical approaches broadly agree in assessing the primary role of in-situ growth for galaxies below \ms$\lesssim 10^{11}\ M_\odot$. However, models become generally more discordant when predicting the evolution of the most massive galaxies. One of the main reasons for such discrepancies can be traced back to the growing significance of the systematics associated with observations such as surface brightness variations, estimates of the proper background, cosmic variance, stellar mass estimates, the number of mergers and the initial mass function \citep[][]{VanDokkum2010,Marchesini2009,Behroozi2013,Maraston2013,Bernardi2014,Shankar2014a,Aversa2015,Leauthaud2016,Bernardi2016b}. In particular, \cite{Bernardi2016} have recently shown that even when homogeneous measurements are carried out at different redshifts, a clear understanding of the evolution of the most massive galaxies still remains elusive. The aim of this paper is to set more stringent and secure constraint on the evolution of the most massive, central galaxies in the local universe for \ms \textgreater $3 \times 10^{11}\ M_\odot$ for which data are still incomplete and/or uncertain, especially at high redshifts. In this work, we use a series of observationally-driven models that, by design, rely on very few assumptions and thus provide us with constraints less clouded by more complex modelling. This paper is structured as follows: In section 2 we give an overview of our methodology and describe our sample selection. In section 3 we discuss the constraints we set on the assembly scenario of massive galaxies. In section 4 we investigate the relative importance of in-situ processes and mergers in driving the evolution of massive ETGs in a late assembly scenario using both observationally informed models as well as a full cosmological, semi-empirical models. We adopt a flat $\Lambda CDM$ cosmological with $\Omega_M = 0.3$, $h = 0.7$, $\Omega_B = 0.045$, $\sigma_8 = 0.8$, $d_c^0 = 1.69$, and assume a Chabrier initial mass function \citep[IMF:][]{Chabrier2003}. Throughout this paper, we define the halo mass as \mh$=M_{200c}$, 200 times the critical density at redshift $z$.

In this work, we have set tighter constrains on the assembly and evolution of massive, central galaxies. We utilise a catalogue of dark matter haloes created from the Bolshoi simulation. We populate these haloes with a stellar mass using recent rendition of the stellar mass to halo mass relation by \cite{Kravtsov2014} and \cite{Shankar2014b} at $z=0$ and select haloes with $\log\ ($\ms$)> 11.5\ M_\odot$. We then trace host haloes back to the putative formation epoch, \zf$=2-4$, as inferred from the stellar ages of massive ETGs. At this epoch, we estimate the total mass in baryons within the halo from the baryon fraction. We find that the stellar mass of the ETG in the local universe is comparable to, if not higher than, the total baryonic mass contained within the progenitor halo. From this comparison, we draw the following important conclusions. \begin{enumerate} \item In-situ formation: For these massive galaxies to have fully assembled at the formation epoch, the efficiency of converting baryons into stars needs to be extremely high if not 100\%. We also show that this assembly scenario would lead to all ETGs being extreme outliers with respect to what is predicted by abundance matching at \zf. \item Size: Even when assuming an extremely efficient star formation at \zf, the galaxy would not be allowed any size growth since the formation epoch. Even an in-situ expansion would in fact require a mass loss of $\geq 70\%$ of the initial baryon content to be sufficiently efficient. Thus, in a strictly monolithic scenario, progenitors of massive galaxies should already be extended systems at their formation epoch. Measurements of the structure of massive galaxies in massive haloes will be critical to assess this possibility. \item Late assembly: Star formation could contribute to the stellar mass growth of the progenitors of massive galaxies, but cannot explain their full evolution. We show through state-of-the-art, cosmological, semi-empirical models that massive galaxies could have indeed assembled most of their final mass via late mergers and be consistent with available data on their size evolution. It remains to be seen the impact of mergers on other (tight) galaxy scaling relations involving velocity dispersion \citep[e.g.,][and references therein]{Bernardi2011b, Bernardi2011a,Shankar2016}. \end{enumerate} More secure and statistically relevant measurements of the stellar mass and structure of high redshift brightest cluster galaxies will be of key relevance to discern between merger scenarios and extremely efficient starbursts events.

1607.05732

1607

1607.02975_arXiv.txt

{} {We aim to study the 250 $\mu$m luminosity function (LF) down to much fainter luminosities than achieved by previous efforts.} {We developed a modified stacking method to reconstruct the 250 $\mu$m LF using optically selected galaxies from the SDSS survey and {\it Herschel} maps of the GAMA equatorial fields and Stripe 82. Our stacking method not only recovers the mean 250 $\mu$m luminosities of galaxies that are too faint to be individually detected, but also their underlying distribution functions. } {We find very good agreement with previous measurements in the overlapping luminosity range. More importantly, we are able to derive the LF down to much fainter luminosities ($\sim25$ times fainter) than achieved by previous studies. We find strong positive luminosity evolution $L^*_{250}(z)\propto(1+z)^{4.89\pm1.07}$ and moderate negative density evolution $\Phi^*_{250}(z)\propto(1+z)^{-1.02\pm0.54}$ over the redshift range $0.02 < z< 0.5$.} {}

Luminosity functions (LF) are fundamental properties of the observed galaxy populations that provide important constraints on models of galaxy formation and evolution (e.g. Lacey et al. 2015; Schaye et al. 2015). Studying the LF at far-infrared (FIR) and sub-millimetre (sub-mm) wavelengths is critical. Half of the energy ever emitted by galaxies has been absorbed by dust and re-radiated in the FIR and sub-mm (Hauser \& Dwek 2001; Dole et al. 2006). The spectra of most IR luminous galaxies peak in the FIR and sub-mm (Symeonidis et al. 2013; Casey et al. 2014). Finally, our knowledge of the FIR and sub-mm LF is relatively poor. The first 250 $\mu$m LF measurement was made by Eales et al. (2009) with observations conducted using the Balloon-borne Large Aperture Submm Telescope (BLAST; Devlin et al. 2009). {\it Herschel} (Pilbratt et al. 2010) significantly improved over BLAST with increased sensitivity, higher resolution, and larger areal coverage. Dye et al. (2010) detected strong evolution in the 250 $\mu$m LF out to $z\sim0.5$, using the {\it Herschel}-Astrophysical Terahertz Large Area Survey (H-ATLAS; Eales et al. 2010). Using the {\it Herschel} Multi-tiered Extragalactic Survey (HerMES; Oliver et al. 2012), Vaccari et al. (2010) presented the first constraints on the 250, 350, and 500 $\mu$m as well as the infrared bolometric (8-1000 $\mu$m) LF at $z<0.2$. More recently, combining {\it Herschel} data with multi-wavelength datasets, Marchetti et al. (2016) derived the LF at 250, 350, and 500 $\mu$m as well as the bolometric LF over $0.02<z<0.5$. Evolution in luminosity ($L^*_{250}\propto(1+z)^{5.3\pm0.2}$) and density ($\Phi^*_{250}\propto(1+z)^{-0.6\pm0.4}$) are found at $z<0.2$. Marchetti et al. (2016), however, were unable to constrain evolution beyond $z\sim0.2,$ as only the brightest galaxies can be individually detected at higher redshifts. Despite the significant progress made, the determination of the LF is still hampered by many difficulties. Large samples over large areas are required for accuracy. We need to focus on smaller areas with increased sensitivity, however, to probe the faint end. At the {\it Herschel}-SPIRE (Griffin et al. 2010) wavelengths, confusion (related to the relatively poor angular resolution) is a serious challenge for source extraction, flux estimation, and cross-identification with sources detected at other wavelengths. In addition, issues such as completeness and selection effects due to the combination of several surveys are extremely difficult to quantify (e.g. Casey et al. 2012). In this paper, we present a new analysis of the 250 $\mu$m LF by stacking deep optically selected galaxy catalogues from the Sloan Digital Sky Survey (SDSS) on the SPIRE 250 $\mu$m images. We bypass some major difficulties in previous measurements (e.g. complicated selection effects, reliability of the cross-identification). The paper is organised as follows. In Section 2, we describe the relevant data products from the SDSS and {\it Herschel} surveys. In Section 3, we explain our stacking method, which recovers the mean properties and underlying distribution functions. In Section 4, we present our results and compare with previous measurements. Finally, we give conclusions in Section 5. We assume $\Omega_m=0.25$, $\Omega_{\Lambda}=0.75$, and $H_0=73$ km s$^{-1}$ Mpc$^{-1}$. Flux densities are corrected for Galactic extinction (Schlegel, Finkbeiner \& Davis 1998).

We study the low-redshift, rest-frame 250 $\mu$m LF using stacking of deep optically selected galaxies from the SDSS survey on the {\it Herschel}-SPIRE maps of the GAMA fields and the Stripe 82 area. Our method not only recovers the mean 250 $\mu$m luminosities $L_{250}$ of galaxies that are too faint to be individually detected, but also their underlying distribution functions.We find very good agreement with previous measurements. More importantly, our stacking method probes the LF down to much fainter luminosities ($\sim25$ times fainter) than achieved by previous efforts. We find strong positive luminosity evolution $L^*_{250}(z)\propto(1+z)^{4.89\pm1.07}$ and moderate negative density evolution $\Phi^*_{250}(z)\propto(1+z)^{-1.02\pm0.54}$ at $z< 0.5$. Our method bypasses some major difficulties in previous studies, however, it critically relies on the input photometric redshift catalogue. Therefore, issues such as photometric redshift bias and accuracy would have an impact. Over the coming years, our stacking method of reconstructing the LF will deliver even more accurate results and also extend to even fainter luminosities and higher redshifts. This is because, although we are probably not going to have any FIR/sub-mm imaging facility that will surpass {\it Herschel} in terms of areal coverage, sensitivity, and resolution in the near future, our knowledge of the optical and near-IR Universe will increase dramatically with ongoing and planned surveys such as DES and LSST. In addition, large and deep spectroscopic surveys such as EUCLID and DESI will further improve the quality of photometric redshift.

1607.02975

1607

1607.02144_arXiv.txt

We investigate systematically whether accreting black hole systems are likely to reach global alignment of the black hole spin and its accretion disc with the binary plane. In low--mass X--ray binaries (LMXBs) there is only a modest tendency to reach such global alignment, and it is difficult to achieve fully: except for special initial conditions we expect misalignment of the spin and orbital planes by $\sim 1$~radian for most of the LMXB lifetime. The same is expected in high--mass X--ray binaries (HMXBs). A fairly close approach to global alignment is likely in most stellar--mass ultraluminous X--ray binary systems (ULXs) where the companion star fills its Roche lobe and transfers on a thermal timescale to a black hole of lower mass. These systems are unlikely to show orbital eclipses, as their emission cones are close to the hole's spin axis. This offers a potential observational test, as models for ULXs invoking intermediate--mass black holes do predict eclipses for ensembles of $\gtrsim 10$ systems. Recent observational work shows that eclipses are either absent or extremely rare in ULXs, supporting the picture that most ULXs are stellar-mass binaries with companion stars more massive than the accretor.

\label{intro} Accreting black holes frequently have their spins at least initially misaligned from the angular momentum of the mass reservoir feeding them. This is generic for supermassive black holes (SMBH) \citep[cf][]{King:2006aa} and is possible in stellar--mass binary systems, particularly if they have undergone a supernova explosion. But any misalignment must evolve as accretion begins. The differential Lense--Thirring precession of disc orbits produces viscous torques on the accretion disc which try to make everything axisymmetric. In stellar--mass binaries, the flux of mass from the companion with angular momentum parallel to the binary axis usually overwhelms these torques in the outer disc, which stays in the binary plane as a result. But close to the black hole the Lense--Thirring effect generally wins, and the inner disc plane rapidly co-- or counter-- aligns with the spin plane on the local precession time \citep{Scheuer:1996aa,King:2005aa}. The transition between the outer disc, aligned with the binary orbit, and the inner disc, aligned with the hole spin, occurs either in a smooth warp \citep{Bardeen:1975aa} or (for larger misalignments) an abrupt break \citep{Nixon:2012aa,Nixon:2012ad,King:2013aa}. We call this configuration -- hole spin and inner disc aligned, but both misaligned from the binary orbit -- {\it central alignment}. We expect this kind of alignment for most discs around compact objects because the the Lense--Thirring effect establishes it very quickly in the inner disc after accretion on to the black hole begins, or resumes. This is also likely for SMBH in active galactic nuclei \citep{King:2006aa,King:2007ab}. But if accretion from a binary companion continues for an extended time, the system may also tend towards a state of {\it global alignment}, where spin, disc and orbital rotation are all parallel or (possibly for the spin) antiparallel. The relative orientation of the hole's spin and the binary axis has a significant effect on the observable properties of accreting stellar--mass black--hole binary systems, so the question of how close a system is to global alignment is important. Studies of it so far either consider individual systems \citep{Martin:2008aa, Maccarone:2002aa,Maccarone:2015aa} or the effect on one method of trying to measure black hole spin \citep{Steiner:2012aa}, which assumes that candidate systems are close to global alignment. Our aim here is to give a systematic picture of whether various types of accreting binaries approach global alignment, including whether this is expected in various models of ultraluminous X--ray sources (ULXs).

We have investigated whether various accreting black hole systems are likely to reach global alignment of the black hole spin and its accretion disc with the binary plane. A fairly close approach to this state is likely in systems where the companion star fills its Roche lobe and transfers mass to a lower--mass black hole. Such systems are promising candidates for ULXs, and are unlikely to show eclipses as their emission cones are close to the hole's spin axis. This offers a potential observational test, as models for ULXs invoking accretion from stellar--mass companions on to intermediate--mass black holes do predict eclipses for an ensemble of $\gtrsim 10$ systems. \cite{Middleton:2016aa} recently showed that eclipses are either absent or extremely rare in among all ULXs for which variability has been measured, in agreement with our result that stellar-mass ULXs should not eclipse because they are close to global alignment. In standard low--mass X--ray binaries there is a modest tendency to reach global alignment, so except for special initial conditions (such as initially retrograde black hole spin) we would expect a misalignment of the spin and orbital planes $\gtrsim 1$ radian. This agrees with the conclusions of \cite{Maccarone:2002aa,Maccarone:2015aa}, and weakens those of \cite{Steiner:2012aa}. It increases the systematic error in attempting to measure black hole spin by comparing the area of the event horizon with that expected from the measured black hole mass. Finally, in high--mass X--ray binary systems, neither spinup nor global alignment is likely.

1607.02144

1607

1607.08250_arXiv.txt

We present a simple model for low-mass planet formation and subsequent evolution within ``transition'' discs. We demonstrate quantitatively that the predicted and observed structure of such discs are prime birthsites of planets. Planet formation is likely to proceed through pebble accretion, should a planetary embryo ($M\gtrsim 10^{-4}$~M$_\oplus$) form. Efficient pebble accretion is likely to be unavoidable in transition disc dust traps, as the size of the dust particles required for pebble accretion are those which are most efficiently trapped in the transition disc dust trap. Rapid pebble accretion within the dust trap gives rise, not only to low-mass planets, but to a large accretion luminosity. This accretion luminosity is sufficient to heat the disc outside the gravitational influence of the planet and makes the disc locally baroclinic, and a source of vorticity. Using numerical simulations we demonstrate that this source of vorticity can lead to the growth of a single large scale vortex in $\sim 100$ orbits, which is capable of trapping particles. Finally, we suggest an evolutionary cycle: planet formation proceeds through pebble accretion, followed by vortex formation and particle trapping in the vortex quenching the planetary accretion and thus removing the vorticity source. After the vortex is destroyed the process can begin anew. This means transition discs should present with large scale vortices for a significant fraction of their lifetimes and remnant planets at large ~10AU radii should be a common outcome of this cycle.

Despite observational advances in characterising exoplanets and the environments in which they form (protoplanetary discs), a coherent picture that explains the origin and diversity of planets from their disky forebears remains elusive. We know planets must sculpt the discs in which they form, and indeed we see structures in observed protoplanetary discs; however, a lack of understanding means that using these observed structures to learn about the details planet formation is challenging. There is a handful of disc-systems which show IR colours and SEDs that are inconsistent with primordial discs: discs that are optically thick from small to large radii \citep[e.g.][]{Strom1989,Skrutskie1990}. Instead, these objects show inner regions that are heavily depleted in dust, while returning to primordial levels in the outer regions \citep[e.g.][]{Calvet2005,Espaillat2014,Owen2016}. Given these discs are partially cleared (in at least the dust) they have been termed ``transition'' discs. It is now known that transition discs do not represent a homogeneous class \citep[e.g.][]{OC12}. While many transition discs are believed to be protoplanetary discs caught in the act of clearing through photoevaporation \citep{Cieza2008,Merin2010,Owen2011}, a significant subset do not match any of the characteristics that would be associated with a protoplanetary disc in the throes of destruction \citep[e.g.][]{Kim2009,Espaillat2010,OC12}. {\bc Approximately half} of transition discs {\bc surveyed by \citet{OC12} had} large accretion rates ($\dot{M}\gtrsim10^{-9}$ M$_\odot$~yr$^{-1}$), large cleared dust cavities (out to $\gtrsim 10$~AU) and they {\bc were} often the brightest of all class II discs at mm wavelengths \citep{Andrews2011,OC12,Ansdell2016}; for this reason they have been termed ``mm-bright transition discs" \citep{Owen2016}. The majority of mechanisms invoked to create a transition disc signature do so by removing dust from the inner regions, and preventing its resupply, by trapping dust outside some radius in a pressure enabled dust trap. Indeed this is true for the two most commonly invoked mechanisms: photoevaporation \citep[e.g.][]{Clarke2001,Alexander2007,Owen2011,Gorti2015} and gap formation by giant planets \citep[e.g.][]{Calvet2005,Rice2006,Zhu2011,Zhu2012,Owen2014}. Several of the mm-bright transition discs have been imaged at high resolution using {\it ALMA}, and in these high resolution images they show strong axisymmetric {\it and} non-axisymmetric emission features \citep[e.g.][]{Casassus2013,vanderMarel2013,Perez2014,vanderMarel2015,Andrews2016,Canovas2016} and see \citet{Casassus2016} for a recent review. These features have been linked to signs of planet formation, although as yet there is no clear understanding of how to link these observations to theory. A common interpretation of the non-axisymmetric structures is that they caused by large-scale vortices \citep[e.g.][]{vanderMarel2013,Ataiee2013,LyraLin2013}. Since these vortices represent local pressure maxima that orbit the star with roughly the local Keplerian velocity they can efficiently trap dust particles \citep[e.g.][]{Meheut2012,LyraLin2013,Zhu2014}. The strong dust density contrast that can result from dust particle trapping in vortices can lead to strong azimuthal surface brightness differences. The Rossby Wave Instability \citep[RWI][]{Lovelace1999,Li2000} provides a mechanism to generate vortices in astrophysical discs that contain radial structure. Sharp radial features give rise to an extremum in potential vorticity that leads to vortex formation. Non-linear hydrodynamic simulations \citep[e.g.][]{Li2001} show that several small scale vortices grow from the linear instability before they begin to grow and merge resulting in one large scale vortex, which can trap particles \citep[e.g.][]{Lyra2009,Meheut2012,Zhu2014}. Transition discs are believed to contain a cavity edge in the gas, and such drops in gas surface density have recently been observed in CO emission \citep{vanderMarel2015,vanderMarel2016}\footnote{\bc Note: CO is likely to be optically thick, and at this stage it is still difficult to directly transform CO structures into gas structures.} . Such an axisymmetric gas structure is necessary to explain many of the observed features of transition discs, however, it remains to be seen whether the observed density drops are steep enough to trigger the RWI. While massive planets {\bc(of order Jupiter mass and higher)} inserted into discs in simulations are known to produce deep and sharp cavities, sharp enough to trigger the RWI, it is unclear whether this is likely to occur in reality. This is because massive planets are inserted into disc simulations instantaneously (or grown over several orbits), necessarily producing a transient phase with a very sharp cavity which is RWI unstable. In a realistic scenario, however, a planet accretes and grows on a time-scale comparable to the evolution of the protoplanetary disc itself. Viscosity present in the disc and, in fact, vortices created by any planet gap initially, can relax the gas to a much smoother distribution, which would be {\it stable} to the RWI over the long-term. This raises the intriguing question: if a transition disc cavity is stable to the RWI, or only unstable for a short period of time, why do transition discs, that are observable on time-scales $\gtrsim 10^4$ orbits, display non-axisymmetric structures generated by RWI vortices? This question is perhaps the biggest weakness of the RWI mechanism in explaining the observed transition disc structures \citep{Hammer2016}. {\bc The RWI can also be triggered in protoplanetary discs at the edge of the dead-zone \citep[e.g.][]{Lyra2009}, where the change in viscosity leads to a sharp change in the surface density \citep[e.g.][]{Gammie1996}. Recent hydrodynamic simulations have shown that large-scale vortices can also be produced this way \citep[e.g.][]{Regaly2012,Flock2015,Lyra2015,Ruge2016} through the RWI. {\bcc Recent work \citep{Ruge2016,Pinilla2016} has shown that the dead-zone model can produce rings and gaps in scattered-light and mm images, although it remains to be seen if the model can explain the large drops in dust-opacity in the inner disc sufficient to reproduce the SEDs of transition discs.} Finally, vortices can also be formed through baroclinic instabilities \citep[e.g.][]{Klahr2003,Lesur2010,Raettig2015}.} By trapping dust at some radius, while it continues to drift in from larger radii, transition discs will have significant increases in the dust-to-gas ratio in the dust trap, from the standard ISM value of 0.01, to values that can approach unity \citep[e.g.][]{Pinilla2012}. Therefore, it has been suggested that these transition disc dust traps are likely to sites of increased planetesimal and planet formation. Indeed \citet{Lyra2008}, demonstrated that it is possible to get direct collapse to form planetary embryos in RWI generated vortices. One interesting advance in the theory of planet formation is the concept of pebble accretion \citep[e.g.][]{Ormel2010,Johansen2010,Lambrechts2012}. Dust particles that have gas-drag induced stopping times comparable to their orbital times can be rapidly accreted by planetary embryos. This is because gas pressure gradients induced by the planetary embryo's presence cause dust particles to be accreted if they approach within the proto-planet's Hill sphere, decreasing the growth time to $\sim 10^{4}$~years --- significantly faster than the standard planetesimal accretion times. Since pebble accretion is very sensitive to dust particle size --- it is only efficient for dust particles that are close to optimally coupled to the gas --- it can only work where there are large numbers of dust particles close to this size. The dust traps in transition discs {\it inevitably} provide a reservoir of these pebbles, as those dust particles that are most efficiently accreted are also those most efficiently trapped. Here we explore the possibility of low-mass planet formation in the pressure traps of transition discs, and argue that it is likely to naturally arise, while perhaps also explaining a variety of observational signatures that are now commonly associated with transition discs. We are agnostic about the mechanism that creates the transition disc itself, but argue that if low-mass planet formation begins in transition discs as we suggest it is: i) likely to be rapid and ii) if the disc is sufficiently massive, it is able to generate vortices that can lead to non-axisymmetric structures similar to those recently observed. We structure our work as follows: in Section~\ref{sec:mechanism} we discuss the physical picture and motivate the basic principles of our new mechanism. In Section~\ref{sec:sims} we present numerical simulations to look at the non-linear long term evolution. We discuss our results in Section~\ref{sec:results} and summarise in Section~\ref{sec:summary}.

\label{sec:results} We have shown that ``transition'' discs are prime sites for the growth of low-mass planets by pebble accretion. This planet formation scenario has two important -- perhaps mutually exclusive -- implications. Firstly, if the planet were able to accrete pebbles at the standard rate for the entire time planet formation is occurring, then it could quickly deplete the local pebble reservoir in a very rapid $\lesssim 10^{5}$~year time-scale. Secondly, if the accretion rate results in a large accretion luminosity, such that it can heat the disc material outside its Hill radius, then it can lead to large scale vortex formation. The first implication begs an interesting question: if the pebble depletion time-scale is so fast $\lesssim 10^{5}$~years, how is it we see a number of transition discs with large mm-fluxes (and hence large reservoirs of pebbles) that almost certainly have lifetimes $\gtrsim 10^{5}$ years? We hypothesise the answer to this question lies in the second consequence. The observed, probably long lived, ``transition'' discs that have large mm-fluxes \citep[e.g.][]{Andrews2011,OC12,Owen2016}, are those discs which are likely to have the highest pebble accretion rates, and as such the discs most prone to vortex formation. Vortex formation, results in an azimuthal pressure trap that can very efficiently trap pebbles \citep[e.g.][]{Meheut2012,Ataiee2013,Birnstiel2013,Zhu2014}. After all, the same dust particles that are likely to be trapped in the dust trap, are those which will accrete onto the planet, and are also those most likely to be trapped in the vortex. Since the vortex can migrate slightly, and does not necessarily have a pattern speed that is exactly Keplerian (as seen in our simulations), then the pebbles that can accrete on to the planet are trapped in the vortex and can only accrete onto the planet for a small fraction of time, or none if the vortex and planet have migrated apart. Therefore, vortex formation and subsequent particle trapping may be the only way to ensure that these mm-bright transition discs remain long-lived. In fact, we can calculate the critical pebble surface density for a transition disc to form a vortex, and in principle remain long lived. We do this by assuming this threshold occurs when $R_T/R_H>1$. Since the Hill radius scales with planet mass as $M_p^{1/3}$ and $R_T$ scales as $M_p^{17/24}$ (where the solid planet mass-radius relationship can be approximated as $M_p=A_{\rm MR}R_p^4$ in the range 1-10 M$_\oplus$ using the profiles from \citealt{Fortney2007}). Thus, the limiting case when a disc can cross the $R_T/R_H>1$ threshold will be at the highest mass the planet can possibly reach. Therefore, if we approximate the maximum mass a pebble accreting planet can reach as $M_p\approx\Sigma_{\rm peb}/2\pi aH_p$ (i.e. the mass it would reach if it had accreted all the pebbles) and assume it is accreting from a reservoir of pebbles with surface density $\Sigma_{\rm peb}$, then we estimate the critical pebble surface density for vortex formation to occur before the entire reservoir is sequestered into a planet as: \begin{equation} \Sigma^{\rm crit}_{\rm peb}=\left[\frac{16\pi\sigma T_d^4A_{\rm MR}^{1/4}}{2G\Omega\left(2\pi a H_p\right)^{3/4}}\right]^{4/7}\label{eqn:sigma_crit} \end{equation} Given a $T\propto R^{-1/2}$ temperature profile for the passively heated, flared disc \citep[e.g.][]{Kenyon1987}, Equation~\ref{eqn:sigma_crit} only depends on a few parameters such that: \begin{equation} \Sigma^{\rm crit}_{\rm peb}\approx 0.3\, {\rm g\, cm^{-2}} \left(\frac{H_p}{H}\right)^{-3/7}\left(\frac{a}{20\,{\rm AU}}\right)^{-5/7}\left(\frac{M_*}{1\,{\rm M}_\odot}\right)^{-4/7}\label{eqn:sigma_crit2} \end{equation} where we have left $H_p/H$ as a free parameter here, but we suspect it to close to unity as discussed above. This critical surface density threshold can be compared to many of the well known mm bright transition discs. This comparison is shown in Figure~\ref{fig:Sigma_crit_compare}, where we plot the peak surface density at the edge of the cavity determined from mm imaging by \citet{Andrews2011} (using their model fits), compared to the result from Equation~\ref{eqn:sigma_crit2}. {\bc This is done by taking the surface density models provided in \citet{Andrews2011}, which are obtained from fits to the mm image and spectral energy distribution. The pebble surface density is then taken to be the surface density in mm-sized particles at the peak of the profile. As such the values are uncertain due to several factors: (i) the simple surface density profile assumed by \citet{Andrews2011}; (ii) uncertainties in the underlying dust-particle distribution which could contribute to the mm-flux and (iii) the trap could be optically thick. } \begin{figure} \centering \includegraphics[width=\columnwidth]{Threshold_compare_2} \caption{The peak surface density of mm particles determined from mm imaging for well known mm-bright ``transition'' discs taken from \citet{Andrews2011} shown as points. The point sizes indicate the stellar mass. The lines show the minimum pebble surface density for vortex formation to be possible indicated by Equation~\ref{eqn:sigma_crit2}, the lines show different stellar masses: 0.5 (dot-dashed), 1.0 (solid) \& 2.0 M$_\odot$ (dashed). {\bc The labels show the individual source names.} }\label{fig:Sigma_crit_compare} \end{figure} Figure~\ref{fig:Sigma_crit_compare} shows that the vast majority of mm-bright ``transition'' discs have mm size particle surface densities sufficiently high that vortex formation due to low-mass planet formation is possible. Finally, it is well known that high viscosities can prevent vortex formation in the planet induced RWI mechanism and can also dissipate vortices \citep[e.g.][]{ValBorro2007,Fu2014b,ZhuStone14}. Our results are also consistent with the previous RWI results, in that large scale vortex formation is suppressed when the typical viscous alpha parameter is $>10^{-3}$. While typical values to explain the global evolution of the protoplanetary discs are of this order \citep[e.g.][]{Hartmann1998,Owen2011} there is every reason to expect the viscosity is likely to be lower in dust-traps and the outer regions of protoplanetary discs. Firstly, non-ideal MHD effects are important in the outer regions of protoplanetary discs \citep[e.g.][]{Armitage2011} and in the case of ambipolar diffusion dominated discs, vortices are known to form and survive for an observable length of time \citep{ZhuStone14}. Secondly, the enhanced dust content in the dust trap is known to suppress the strength of MRI turbulence \citep[e.g.][]{Jacquet2012}. Therefore, it is not unreasonable to assume that the viscosity in the neighbourhood of the dust-trap is sufficiently low to allow the formation of vortices. \subsection{Long term evolution} The long term evolution of the disc will be strongly controlled by whether it is able to form a vortex or not. As discussed above, if a large scale vortex is able to form it will trap all the pebbles in the vortex itself. As none of the vortices seen in the simulations are co-located with the orbiting planet \citep[see][for a discussion of how the planet and vortex may interact]{Ataiee2014}, then the vortex will starve the planet of pebbles and the rapid pebble accretion will cease. Our simulations suggested that vortex formation is fairly rapid $\sim 100$ orbits. This is similar to the dust trapping time-scale, thus we suspect that the pebbles will become easily trapped in the vortex. At this stage the pebbles will no longer be available to rapidly accrete onto the planet. If the planet's accretion source is quickly shut-off then it will still be luminous for a short period of time, as the gaseous envelope previously supported by the accretion contracts towards the planet. This contraction will maintain its luminosity for a time-scale roughly similar to the gaseous envelopes Kelvin-Helmholtz time-scale ($t_{\rm KH}$), which for a planet with an envelope mass considerably less than the solid mass is given by, \citep[e.g.][]{Lee2015}: \begin{equation} t_{\rm KH}=\frac{GM_p^2X_{\rm env}}{R_pL}=X_{\rm env}t_{\rm acc} \end{equation} Thus, the Kelvin-Helmhotz time-scale is shorter than the accretion time-scale by the gas envelope to planet mass ratio. For these masses and luminosities $X_{\rm env}$ will be in the range $\sim$0.01, indicating that the luminosity output of the planet will drop on a time-scale $\sim 10-100$ orbits. This means that once the large scale vortex has trapped the majority of pebbles, the source of the barcolinicity generating the vortex will disappear on a comparable time-scale to the vortex formation time-scale, approximately as $R_T\propto t_{\rm KH}^{-1/2}$. Since a gravitationally contracting object will follow an evolution such that $t_{\rm KH}$ is approximately the time it has been cooling, then $R_T$ will decrease as the square root of time. Now, since the vortex lifetime is finite, as without a source of barcolinicty, viscosity, instabilities \citep[e.g.][]{Lesur2009} or even the inertia of the dust itself \citep[e.g.][]{Fu2014} can destroy a large scale vortex on a time-scale of roughly 1000s of orbits, which is short compared to the disc's lifetime (e.g. $> 10^4$ orbits). We can therefore expect that vortex generation and dissipation occurs several times during the disc's lifetime, where pebble accretion onto a planet generates a vortex, which traps the pebbles hence suppressing the accretion and vortex generation. The vortex then dissipates after some time-scale, releasing the dust particles into an axi-symmetric ring that can then undergo pebble accretion onto a planet again, forming another vortex. The entire process can repeat for a long time-scale until the pebble reservoir is too heavily depleted to permit vortex formation. If the pebble reservoir is too small to permit vortex formation, {\bc although as demonstrated by Equation~\ref{eqn:sigma_crit2} still significant to make the disc appear as a mm-bright disc.} As discussed in Section~2, pebble accretion onto a planet will then rapidly deplete its local reservoir on a time-scale $\lesssim 10^{5}$~years. {\bc This means the time-scale that a disc in this model would appear at intermediate mm-fluxes would be short.} Once the planet has depleted its local reservoir, particles at larger orbital separations will continue to drift into the dust trap. At large radii the maximum dust particle size is limited by drag rather than fragmentation \citep[e.g.][]{Birnstiel2012b}. Therefore, we can imagine that as dust particles drift from large radii in the disc, they will begin to grow and when they arrive in the vicinity of the planet they will be accreted readily as they will have naturally reached a size with $\tau_s=1$ at the planet's radius. In this stage, the planet's accretion rate will be limited by the accretion rate of dust particles into the trap due to drift, rather than the standard pebble accretion rate. Since these planets will necessary be low-mass ($< 10$~M$_\oplus$) they are unlikely to be able to accrete a significant gas envelope \citep[e.g.][]{Rafikov2006,Piso2014}. Thus, unlike the dust particles that could in principle be rapidly sequestered into low-mass planets, the gas will remain largely unchanged. This means this process may ultimately result in a low-mass dust-poor, gas-rich, low-mass-planet-rich disc, a type of disc for which there are poor observational constraints. Finally, the last concern is migration. These planets are low enough mass that they are unlikely to open a gap in the gas disc, where the gap opening mass is more typically $\sim 0.1$M$_J$ at tens of AU \citep[e.g.][]{Crida2006}. Therefore, migration is likely to take place in the type-I regime. While the migration rates for low-mass planets are still greatly uncertain in realistic protoplanetary discs, even the isothermal type-I migration rates, which are likely to strongly over-predict the migration rates, give rates $\sim 10^5$~years for a few earth mass planet at 10s of AU \citep[e.g.][]{Ward1997}. This means that a forming planet is able to initiate vortex formation and it will not migrate away in on the time scale for vortex formation to occur. However, migration might push the formed planet inside the transition disc cavity after the vortex has formed but before it has dissipated. The vortex will certainly qualitatively effect the migration of such a low-mass planet (we note again we do not consider migration of our planet in our simulations presented here), as demonstrated by \citet{Ataiee2014}. This might mean that over a Myr lifetime of planet formation, the cycles of vortex generation and dissipation in the disc may dump a handful of low-mass planets inside the transition disc cavity, which could subsequently scatter. \subsection{Observational implications} While we have suggested that transition discs are prime sites for planet formation through pebble accretion, the fact that it is so rapid means that if it was left to proceed as normal in discs with the parameters of standard mm-bright transition discs ($a\sim 10-50$~AU, $\Sigma_{\rm peb}\sim 1-10$g cm$^{-2}$), it would rapidly deplete the entire disc on an incredible short time $\lesssim 10^{5}$~years. We suggest the fact that these mm-bright transition discs can exist for a long enough time-scale to be observable, is that at the expected dust surface densities in these discs, the act of low-mass planet formation in their dust traps should result in large-scale vortex formation. As discussed above these dust traps should trap all the pebbles within the vortex itself and prevent further planet growth. Therefore, many of these discs should spend {\bc some} fraction of their lifetimes with large scale vortices, which will be observable as large scale asymmetries when the discs are imaged at mm wavelengths. {\bcc For the mm-bright transition discs imaged at high resolution IRS48 \citet{vanderMarel2013} and HD142527 \cite{Casassus2013} show strong asymmetries, and LkH$\alpha$330 \citep{Isella2013} and SAO2016462 \citep{Perez2014} show weaker asymmetries. Many other observed transition discs do not show any evidence of an asymmetry (e.g. LkCa15, SR24S \citealt{vanderMarel2015}, Sz91 \citealt{Canovas2016}, DoAr 44 \citealt{vanderMarel2016}). In recent ALMA surveys of protoplanetary discs \citep[e.g.][]{Pascucci2016,Ansdell2016} several more transition disc like structures have been detected, many of which show no evidence for an asymmetry at $\sim 0.3"$ resolution. To date it is unclear exactly what fraction of transition discs show asymmetries. The reason for this is partly due to unsystematic observations at a variety of resolutions and partly due to the fluid definition of a ``transition disc'': for example should discs with large mm-holes but primordial SEDs be counted in this sample? \citep[See][for a discussion of these discs.]{Andrews2011,Owen2016}. What can be said with any confidence is that a moderate fraction of mm-bright transition discs show asymmetries.} Finally, exactly what kind of asymmetry our vortices produce will depend on the observational wavelength, surface density distribution and other factors (e.g. dust-growth and destruction) we do not consider here. Therefore, we must wait until dust-gas simulations are performed, before we can draw any hard conclusions from the (incomplete) transition disc statistics we have today. We suspect the hot-spot generated by the planet will fade on a time-scale of $\sim 100$ orbits. Thus if the vortex lifetime is significantly longer than 100 orbits, it would be very unlikely to directly observe the effect of the increase in local disc temperature in the continuum image as $R_T$ will have contracted to well inside the planet's Hill sphere once the vortex had trapped all the dust particles. However, it may be possible that the hot spot would generate some chemical fingerprint, that would last longer than the hot spot, spread into a ring, and possibly be a signature of this process. The critical pebble surface density for vortex formation derived in Equation~\ref{eqn:sigma_crit2} is similar to the value required to give a mm-flux at 140pc of $\sim 30$ mJy \citep[e.g.][]{Andrews2005}. This mm-flux value is the discriminating value between mm-bright and mm-faint transition discs determined by \citet{OC12}. \citet{OC12} and \citet{Owen2016} argued that mm-bright transition discs are likely to be rare and long lived (lifetimes $> 10^{5}-10^{6}$ years), whereas mm-faint transition discs are thought to be common (in the sense that all discs experience this phase) and short lived with lifetimes $\lesssim 10^{5}$~years. Here we suggest that this critical pebble surface density for vortex formation could provide the link as to why mm-bright transition discs are likely long lived and mm-faint transition discs are likely short lived. Since mm-bright transition discs can form vortices, they can prevent the dust from being rapidly turned into planets, producing many cycles of planet formation, vortex formation and subsequent destruction. Finally, one obvious consequence of this mechanism is the production of low-mass ($\lesssim$ Neptune mass) planets with orbital separations of tens of AU. There is currently limited observational sensitivity to the low-mass exoplanet population at large separations. However, mirco-lensing surveys have detected Neptune mass planets at separations $\sim$10~AU \citep[e.g.][]{Gaudi2012, Shvartzvald2016} and have suggested that such planets are common \citep[e.g.][]{Gould2006}. The sensitivity of these experiments will only improve with future surveys hopefully revealing the full mass-spectrum and semi-major axis distribution of these systems. \subsection{Future Directions} In this work we have mainly argued that transition disc dust traps are likely to be prime sites of planet formation by pebble accretion. The rapid nature of the planetary accretion is likely to have many interesting consequences. Here we have argued that by modifying the temperature structure outside the gravitational influence of the planet it can provide a source of vorticity, which allows the growth of large scale vortices. In the simulations performed in this work we have described a very ideal setup, where we impose a local hot-spot that is not coupled to the subsequent dynamics. Specifically, we do not to attempt to model the coupled evolution of the dust, gas and planetary accretion in a self consistent way. While the time-scales suggest that large scale vortex formation is likely to occur the finer details of the results require further investigation. The vortex strength, how it traps dust particles, and when and how accretion onto the planet is shut-off due to the fact the pebbles are now trapped in the vortex requires coupled dust and gas simulations, possibly including grain growth. Such simulations will be able to investigate the cycle of planet formation, vortex growth, planet migration and dissipation that allows these discs to exist for the $>10^5-10^6$~years that we have postulated above. Furthermore, in-order to isolate the physics of the problem at hand we have neglected the self-gravity of the disc and the indirect potential that could effect the planetary dynamics and migration. In massive discs, vortex formation could result in the vortices transitioning to global ``fast'' modes \citep{Mittal2015}, similar to the transition of RWI generated vortices into fast modes discussed by \citep{Zhu16a,Zhu16b}. It maybe interesting to investigate the interaction between pebble accretion generated vortices and the indirect potential to see if fast modes can be triggered in this case, even without the original transition disc cavity. Finally, we note that while we have primarily focused on pebble accretion and vortex generation in transition discs, the mechanism of vorticity generation is not limited specifically to transition discs. Indeed, pebble accretion has been invoked to solve numerous planet formation problems at various locations and rates throughout primordial protoplanetary discs \citep[e.g.][]{Bitsch2015}. In our work, the gas cavity prevents the vortices from rapidly migrating allowing them to grow to large sizes. In a primordial disc, we speculate that smaller vortices may be generated which can still trap particles and migrate. By trapping particles this could affect the assumed pebble accretion rates in such calculations. This may be particularly important at smaller radii, as our critical vortex formation threshold scales as $R^{-5/7}$, whereas disc surface densities are thought to fall in a steeper manner with radius, with mm observations of protoplanetary discs suggesting an $R^{-1}$ \citep{Andrews2009} decline, and the Minimum Mass Solar Nebula (MMSN) scaling as $R^{-3/2}$. Therefore, investigating the prospects of vortex generation in a primordial disc that is forming planets through pebble accretion is certainly a worthwhile investigation.

1607.08250

1607

1607.07918.txt

We present 31.5 \um~imaging photometry of 11 nearby Seyfert galaxies observed from the Stratospheric Observatory For Infrared Astronomy (SOFIA) using the Faint Object infraRed CAmera for the SOFIA Telescope (FORCAST). We tentatively detect extended 31 \um~emission for the first time in our sample. In combination with this new data set, subarcsecond resolution $1-18$ \um~imaging and $7.5-13$ \um~spectroscopic observations were used to compute the nuclear spectral energy distribution (SED) of each galaxy. We found that the turnover of the torus emission does not occur at wavelengths $\leq$31.5 \um, which we interpret as a lower-limit for the wavelength of peak emission. We used \textsc{Clumpy} torus models to fit the nuclear infrared (IR) SED and infer trends in the physical parameters of the AGN torus for the galaxies in the sample. Including the 31.5 \um~nuclear flux in the SED 1) reduces the number of clumpy torus models compatible with the data, and 2) modifies the model output for the outer radial extent of the torus for 10 of the 11 objects. Specifically, six (60\%) objects show a decrease in radial extent while four (40\%) show an increase. We find torus outer radii ranging from $<$1pc to 8.4 pc.

The unified model \citep{A1993,UP1995} of active galactic nuclei (AGN) posits that all AGN are essentially the same type of object viewed from different lines of sight. This orientation-based model depends on a circumnuclear toroidal region of optically and geometrically thick dust which can obscure a central region of high-energy emission (a super massive black hole of $\sim$10$^{6-9} M_{\odot}$ and accretion disk), responsible for producing high energy photons. Observed broad (FWHM $\sim$ 10$^{3}$ - 10$^{4}$ km/s) and narrow (FWHM $<$ 10$^{3}$ km/s) emission line features in Type 1 AGN are due to a direct view of the central engine, whereas the circumnuclear dust torus obscures a direct view of the central engine and broad line emission region in Type 2 AGN. Strong support for the unification scheme came from spectropolarimetric observations of NGC 1068 \citep{AM1985}, a Type 2 Seyfert galaxy. It was shown that NGC 1068 contains polarized broad emission lines that are hidden from direct view, but clearly revealed in polarized radiation, proving that this Seyfert 2 has properties similar to a Seyfert 1. Subsequently, evidence for a hidden broad line region was also found in several other highly polarized Seyfert 2's \citep[e.g.][]{MillGood,Brindle1990,Tran}. In both Seyfert 1 and 2 galaxies, dust grains in the torus absorb optical and ultraviolet radiation from the central engine and re-radiate in the infrared (IR). It was assumed early on that most of the dust in the torus must have been distributed in molecular dust clouds or would otherwise not survive expected temperatures \citep{KB1988}; hence early models assumed a homogeneous distribution of dust \citep{PK1992,GD1994,EfRR1995,Granato1997,Sieb2004} for its computational simplicity. In the case of homogeneous models, the amount of mid-IR (MIR) to far-IR (FIR) emission suggested a torus outer extent of up to a few hundred parsecs. These models were based on moderate spatial resolution ($\gg$1 arcsec) photometry, and thus suffered from contamination from diffuse dust emission and stellar emission in the core of the host galaxy. High spatial resolution MIR imaging observations on 8-m class telescopes quickly ruled out a torus of such large extent. Specifically, N- and Q-band observations from Gemini South provided an upper limit on the outer radius of $\sim$2 pc for Circinus \citep{Packham2005} and 1.6 pc for Centaurus A \citep{Radomski2008}. Further N-band interferometric observations on VLTI/MIDI \citep{Jaffe2004,Tristram2007,Raban2009,B2013} and more recently sub-millimeter observations from ALMA \citep{GB2016} confirmed a smaller radius for several objects. The small size of the torus is effectively modeled by an inhomogeneous, ``clumpy" dust distribution throughout its volume. In this scenario, dust of differing temperatures can exist at the same radius \citep{Nenkova2002,Schart2005}, i.e. the illuminated face of one cloud (emitting in near-IR (NIR)) can exist at the same radius as the shadowed face of another cloud (emitting in MIR). Models assuming a clumpy distribution \citep{Nenkova2002,Nenkova2008,Nenkova2008b,Honig2006,Stalevski2012,Sieb2015} also account for the variety of spectral energy distributions \citep[SEDs;][]{Fadda,AH2003,RA09,RA11,AH2011,Lira2013,Ichi2015} that are seen. For example, the NIR to MIR SED is sensitive to the inclination angle of the torus and its width, with the hotter dust within the torus contributing to flatten the SED. The outer radius is best constrained by FIR emission \citep{RA11b,AR2013} since temperatures are generally cooler further away from the central engine, though in detail depends on the total dust distribution. Consequently, knowing the wavelength of peak emission from the torus gives insight as to its radial size. Some models show peak IR emission from the torus between $\sim$$20 - 40$ \um~\citep{Nenkova2008,H2010,Mull2011,Feltre2012}. Thus, observations at wavelengths longer than 20 \um~are essential in determining the wavelength of peak emission. Observing AGN at the highest possible spatial resolution ensures minimal contamination of the signal from the surrounding diffuse dust and stellar emission in the core of the host galaxy. To optimally constrain the torus model parameters, several authors \citep{RA09,H2010,RA11,AH2011,Ichi2015} have combined subarcsecond-resolution $1 - 20$ \um~observations to evaluate the SED. Although these previous studies have effectively described the parameters taking into account high spatial resolution observations at wavelengths $\lambda < 25$ \um, the lack of high spatial resolution observations at wavelengths $\lambda > 25$ \um~leaves the SED at longer wavelengths largely unconstrained, reducing the quality of the fitting to the clumpy models. Unfortunately, since the atmosphere is opaque in FIR, large aperture observations from the ground are impossible at these wavelengths. The space-based \textit{Spitzer} telescope covers this range, but the diameter (0.85m) severely limits its resolution ($\sim$9.1 arcsec at 30 \um). NASA's Stratospheric Observatory For Infrared Astronomy (SOFIA) presents a unique solution for non-space-based observations in the FIR. The 2.5-m telescope on board the aircraft is three times the size of \textit{Spitzer}, providing a much better spatial resolution ($\sim$3.4 arcsec at 30 \um). In this paper, we present new 31.5 \um~imaging data from NASA's SOFIA telescope for a sample of 11 nearby Seyfert galaxies. We used the \textsc{Clumpy} torus models of \citet{Nenkova2008,Nenkova2008b} and a Bayesian approach \citep{AR2009} to fit the IR ($1-31.5$ \um) nuclear SEDs in order to constrain torus model parameters. We aim to determine the AGN contribution to the flux within the SOFIA aperture, estimate the potential turnover of the IR torus emission, and determine the effect on model fits after adding the extended wavelength range. Section \ref{OBS} of this paper contains the observations and data reduction; Section \ref{results} contains our photometric analysis method and its results; Section \ref{MOD} gives the results from the model fitting with the addition of our 31.5 \um~photometric point; Section \ref{DIS} contains an analysis of the torus model parameters; and Section \ref{CON} gives our conclusions. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%% OBSERVATIONS AND DATA REDUCTION %%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \vspace{-0.7cm}

\label{CON} We present new 31.5 \um~imaging data from NASA's SOFIA/FORCAST for 11 nearby Seyfert galaxies. To derive AGN-dominated fluxes within the 3.4 arcsec SOFIA aperture, we used the PSF scaling method of \citet{Radomski2003,RA09,Mason2012}. We then performed a spectral decomposition using the routine of \citet{HC2015} to estimate the contribution of host galaxy components within SOFIA's FWHM and to verify the PSF-scaling results. We also compiled NIR and MIR fluxes from the literature based on previous works \citep{RA09, AH2011, Ichi2015}, and used the \textsc{Clumpy} torus models of \citet{Nenkova2008,Nenkova2008b} togther with a Bayesian approach \citep{AR2009} to fit the IR ($1.2-31.5$ \um) nuclear SEDs in order to constrain torus model parameters. The $1-31.5$ \um~SEDs presented here do not show a turnover below 31.5 \um. The predicted turnover occurs between $30-50$ \um. Further observations in the $32-40$ \um~regime would be beneficial in observationally finding the peak IR emission, thus providing further insight to the torus outer limit. To further investigate this wavelength range, we have been allocated more observation time on SOFIA in 2016 to explore our current AGN sample, as well as 11 more objects, at a wavelength of 37.1 \um. In the next decade, the 6.5-m Tokyo Atacama Observatory (TAO) may also present a unique opportunity to observe AGN in this wavelength range. Its MIMIZUKU instrument will cover $2 - 38$ \um~with a spatial resolution of $1-2$ arcsec at 30 \um~\citep{Kamizuka2014}. By including the 31.5 \um~photometric point in the SED, the model output for the radial extent is modified for 10 of the 11 objects. Six galaxies (60\%) show a decrease in radial extent while four galaxies (40\%) show an increase. The average value from the global posterior of $Y$ decreases from 20 to 17, with an average relative change of $\sim$30\% for the AGN in our sample, supporting the results of \citet{AR2013} which suggest that observations using FORCAST will provide substantial constraints to the \textsc{Clumpy} models. We find torus outer radii in the range $<$1pc to 8.4 pc, consistent with interferometric and high angular resolution MIR observations \citep{Jaffe2004,Packham2005,Tristram2007,Radomski2008} and with recent sub-mm observations from ALMA \citep{GB2016}. \vspace{-0.7cm}

1607.07918

1607

1607.03973_arXiv.txt

The second-order vector mode is inevitably induced from the coupling of first-order scalar modes in cosmological perturbation theory and might hinder a possible detection of primordial gravitational waves from inflation through 21cm lensing observations. Here, we investigate the weak lensing signal in 21cm photons emitted by neutral hydrogen atoms in the dark ages induced by the second-order vector mode by decomposing the deflection angle of the 21cm lensing signal into the gradient and curl modes. The curl mode is a good tracer of the cosmological vector and tensor modes since the scalar mode does not induce the curl one. By comparing angular power spectra of the 21cm lensing curl mode induced by the second-order vector mode and primordial gravitational waves whose amplitude is parametrized by the tensor-to-scalar ratio $r$, we find that the 21cm curl mode from the second-order vector mode dominates over that from primordial gravitational waves on almost all scales if $r \lesssim 10^{-5}$. If we use the multipoles of the power spectrum up to $\ell_{\rm max} = 10^{5}$ and $10^{6}$ in reconstructing the curl mode from 21cm temperature maps, the signal-to-noise ratios of the 21cm curl mode from the second-order vector mode achieve ${\rm S/N} \approx 0.46$ and $73$, respectively. Observation of 21cm radiation is, in principle, a powerful tool to explore not only the tensor mode but also the cosmological vector mode.

Recent development of cosmological observations plays an important role in establishing standard cosmology, namely, the $\Lambda$CDM model. In the current states of observations of cosmological perturbations, for example, we have detected the cosmic microwave background (CMB) temperature anisotropy, the CMB E-mode polarization, galaxy clustering, and so on \cite{Tegmark:2006az,Hinshaw:2012aka,Sanchez:2013tga,Ade:2015xua}. In the context of the cosmological perturbation theory, the perturbations can be decomposed into the scalar, vector, and tensor modes. Current observations show good agreement with perturbations of the first-order scalar mode. The vector and tensor modes are becoming the next observational targets. The vector and tensor modes have been considered in many contexts. The standard model of inflation in the early Universe, namely, the single-field slow-roll inflation, predicts the existence of primordial gravitational waves (PGWs) that correspond to the tensor mode. The vector mode can be induced by cosmological magnetic fields \cite{Lewis:2004ef,Lewis:2004kg}, cosmological defects \cite{Pen:1997ae,Durrer:1998rw,Horiguchi:2015xsa}, or additional vector fields \cite{Zuntz:2010jp,Saga:2013glg}. However, the amplitude of the vector mode generated in these models highly depends on their model parameters. When we expand cosmological perturbations up to the second order, the nonlinear coupling of the first-order scalar modes naturally induces the second-order vector and tensor modes \cite{Assadullahi:2009jc,Ananda:2006af,Saga:2014jca,Baumann:2007zm,Saga:2015apa,Ichiki:2006cd,Fenu:2010kh,Saga:2015bna}. Since the amplitude of the first-order scalar mode is precisely determined by recent observations, the amplitudes of second-order vector and tensor modes can be predicted without introducing additional model parameters. The second-order mode is one of the good observational targets since those modes always exist. Furthermore, the detection of cosmological vector or tensor modes is important to test the validity of the scalar, vector, and tensor decomposition. Weak lensing is one of the tools that can be used to study the cosmological vector or tensor modes. CMB photons are deflected by foreground perturbations such as density fluctuations, the gravitational wave background, or vector perturbations. We can split the deflection pattern into the gradient (parity-even) and curl (parity-odd) modes \cite{Namikawa:2011cs,Yamauchi:2012bc,Yamauchi:2013fra}. The curl mode is induced only from the vector and tensor modes. Therefore, the curl mode is one of the good tracers to explore the cosmological vector or tensor modes. In our previous study \cite{Saga:2015apa}, we discussed the detectability of the second-order curl mode in the CMB lensing and cosmic shear. Unfortunately, because the signal from the second-order curl mode is small, we concluded that we could not detect the second-order curl mode even with an ideal experiment for the full sky without the instrumental noise if we utilize the quadratic estimator method. However, we found that the curl mode from the second-order vector mode is comparable to that from PGWs with tensor-to-scalar ratio $r< 0.1$, especially so in lower redshifts because the second-order vector mode is continuously generated, while PGWs always decay in time. In other words, when there is an observation that enables us to detect PGWs with $r< 0.1$ through lensing, we can also detect the second-order vector mode. In previous studies \cite{Book:2011dz,Masui:2010cz,Sigurdson:2005cp}, it was shown that the 21cm lensing has a possibility, in principle, to detect PGWs with a quite small tensor-to-scalar ratio. Long before reionization begins, no astronomical objects exist, and this era is called the dark ages. Neutral hydrogen atoms emit 21cm line radiation that originates from the hyperfine structure; see, e.g., Ref.~\cite{Furlanetto:2006jb}. In principle, we can observe the 21cm radiation from the redshift $z\approx 200$ to $30$ in future experiments. 21cm photons are deflected by the foreground scalar, vector, and tensor modes. Moreover, we can decompose the deflection angle of the 21cm photons into the gradient and curl modes depending on the parity. Compared with CMB fluctuations, 21cm radiation does not suffer from a diffusion mechanism such as Silk dumping and the 21cm fluctuations on small scales remain until today. Consequently, the available information from 21cm fluctuations is dramatically improved compared with that from CMB fluctuations. Furthermore, 21cm radiation is emitted from each redshift and many maps are available. For the above reason, 21cm lensing reconstruction noise would become quite small compared with CMB lensing reconstruction noise. Therefore, although second-order vector and tensor signals tend to be small, there is a possibility to detect these second-order signals in 21cm lensing. In this paper, we focus on the 21cm lensing curl mode induced from the second-order vector mode. Our aim is to estimate the signal-to-noise ratio of the 21cm curl mode from the second-order vector mode in ideal experiments. In standard cosmology, the first-order vector mode always decays and is neglected in linear theory. The detection of the cosmological vector mode is quite important because it would become a proof of the cosmological perturbation theory itself and the validity of the scalar, vector, and tensor decomposition. This paper is organized as follows. This study is based on the physics of 21cm radiation, cosmological perturbation theory expanded up to second order, and weak lensing. In Sec.~\ref{sec: pre}, we briefly review these physics and their mathematical basics with appropriate references. To predict the 21cm radiation fluctuation, we need to solve the perturbed Boltzmann equation for the 21cm photons. The Boltzmann equation for 21cm photons has the collision term with neutral hydrogen atoms. We solve the Boltzmann equation for 21cm photons numerically. The second-order vector mode is generated from the coupling of the first-order scalar potentials and is solved numerically in a straightforward manner. To discuss the detectability, we derive the noise spectrum induced from the lensing reconstruction in a similar way as in CMB lensing reconstruction. The advantage of 21cm lensing reconstruction is that one can coadd different redshift slices. In Sec.~\ref{sec: results}, we show our main results and give some discussions regarding the 21cm lensing signal and detectability. Finally, we present our conclusions in Sec.~\ref{sec: summary}.

\label{sec: results} In Fig.~\ref{fig: cls}, we show the angular power spectra of the curl mode induced from PGWs with $r = 0.1$ and the second-order vector mode. \begin{figure}[t] \begin{center} \includegraphics[width=0.49\textwidth]{Signal_PGW} \includegraphics[width=0.49\textwidth]{Signal_VEC} \end{center} \caption{% The angular power spectra of the curl mode induced by PGWs with $r = 0.1$ (left) and the second-order vector mode (right) for redshifts from $z = 200$ to $0.6$ as indicated in the figures. The curl mode from PGWs is substantially suppressed on small scales compared with that from the second-order vector mode.} \label{fig: cls} \end{figure} We can see that the lensing signal from PGWs is suppressed as the redshift decreases. On the other hand, the curl mode from the second-order vector mode remains almost constant. We find that the redshift dependence of the second-order vector mode is similar to that of the gradient mode from the first-order scalar potential \cite{Lewis:2006fu}. This is because the second-order vector mode is also sourced from the first-order scalar potential. Therefore the amplitude of the second-order vector mode can have a greater amplitude than the curl mode from PGWs. The amplitude of the curl mode from the second-order vector mode is greater than that from PGWs with $r = 0.1$ on smaller scales, such as $\ell \gtrsim 20$. Furthermore, when the tensor-to-scalar ratio is quite small, e.g., $r \lesssim 10^{-5}$, the curl mode from the second-order vector mode dominates over almost all scales. From this fact, we can conclude that even if we consider ideal observations, it would be difficult to hunt the tensor-to-scalar ratio $r \lesssim 10^{-5}$ by using 21cm lensing. In Fig.~\ref{fig: SN result}, we depict the signal-to-noise ratio for two different values of $\ell_{\rm max} = 10^{5}$ and $10^{6}$, which is our main result in this paper. \begin{figure}[t] \begin{center} \includegraphics[width=0.49\textwidth]{SN_VEC5} \includegraphics[width=0.49\textwidth]{SN_VEC6} \end{center} \caption{% The angular power spectrum of the curl mode from the second-order vector mode and the coadded reconstruction noise by using $\ell_{\rm max} = 10^{5}$ (top left) and $\ell_{\rm max} = 10^{6}$ (top right). Bottom: The signal-to-noise ratio estimated by Eq.~(\ref{eq: SN formula}). For reference, we also show the curl mode signal and signal-to-noise ratio induced by PGWs with $r = 10^{-5}$ and which tensor-to-scalar ratio corresponds to the same amplitude of the second-order curl mode at $\ell = 2$. } \label{fig: SN result} \end{figure} For reference, we also show the signal-to-noise ratio for the case of PGWs with $r=10^{-5}$. PGWs with $r = 10^{-5}$ have almost the same amplitude of the curl mode from the second-order vector mode at $\ell = 2$. In the case of $\ell_{\rm max} = 10^{5}$, the signal-to-noise ratio reaches ${\rm S/N}\approx 0.11$ for the PGWs and ${\rm S/N}\approx 0.46$ for the second-order vector mode and it would be difficult to detect the second-order vector mode and PGWs with $r = 10^{-5}$. On the other hand, in the case of $\ell_{\rm max} = 10^{6}$, we obtain ${\rm S/N}\approx 4.5$ for the PGWs and ${\rm S/N}\approx 73$ for the second-order vector mode. The above signal-to-noise ratio is derived by adopting the reconstruction noise spectrum from the quadratic estimator performed in Refs.~\cite{Namikawa:2011cs,Ade:2015zua}. Reconstruction noise from the quadratic estimator is limited by the cosmic variance of the lensed CMB fluctuations. Ultimately, the iterative estimator proposed in Ref.~\cite{Hirata:2003ka} can reduce the reconstruction noise to zero. Even in that case, the fact that the curl mode from PGW with $r\lesssim10^{-5}$ is concealed by that from the second-order vector mode does not change. The signal-to-noise ratio of the second-order vector mode can reach higher than that of PGWs. PGWs do not induce the curl mode amplitude on smaller scales since PGWs decay on subhorizon scales. On the other hand, the second-order curl mode can remain large on smaller scales and at low redshift since the second-order vector mode is continuously sourced by the first-order scalar gravitational potential. The second-order vector mode grows on subhorizon scales. From this nature, the second-order vector mode may be easier to detect than PGWs on small scales. There is another source of the curl mode, that is, the lens-lens coupling examined in Refs.~\cite{Jain:1999ir,Sarkar:2008ii,Cooray:2002mj}. The lens-lens coupling is sourced by the higher-order Born correction. However, this correction mainly contributes the curl mode on small scales such as $\ell \gg 10$. The curl mode on large scales is important to distinguish the PGWs and the second-order vector mode since the PGWs and the second-order vector mode affect the curl mode on large scales, that is, $\ell \lesssim 10$. When we consider the curl mode on all scales, the lens-lens coupling and the second-order vector mode should be taken into account. To close this section, we describe a feature of the second-order vector mode. The second-order vector mode does not have the free parameter since its source, that is, the first-order scalar mode, is well determined by the current cosmological observations. Therefore, the prediction of the 21cm lensing curl mode from the second-order vector mode is quite robust. In this paper, we studied the detectability of the second-order vector mode by using 21cm radiation from the dark ages. 21cm radiation during the dark ages is a powerful tool to explore small signals such as second-order signatures since 21cm radiation anisotropy on small scales makes multipole moments available up to $\sim 10^{6}$. Furthermore, by multifrequency observations, we can use many redshift slices to decrease the lensing reconstruction noise. We focused on the weak lensing signal of the 21cm radiation from the dark ages. As well as the CMB lensing, the 21cm photons are deflected by the foreground scalar, vector, and tensor perturbations. The deflection angle of the 21cm photons can be decomposed into the scalar (gradient mode) and pseudoscalar (curl mode) potentials depending on its parity. The curl mode is a good tracer of the cosmological vector and tensor modes since the scalar mode induces only the gradient mode. It is known that the second-order tensor mode is the subdominant component in the large-scale structure signal such as weak lensing. On the other hand, the second-order vector mode can have a comparable contribution on large-scale structure to primordial gravitational waves. Accordingly, the observation that can detect PGWs with small tensor-to-scalar ratio can also be used to detect the second-order vector mode with high signal-to-noise ratio. We discussed the detectability of the 21cm lensing curl mode induced from the second-order vector mode for the first time. If the available multipole is limited to $\ell \lesssim 10^{5}$, the 21cm lensing curl mode from the second-order vector mode cannot be detected. If we can extend the maximum multipole up to $\ell_{\rm max} \approx 10^{6}$, the signal-to-noise ratio reaches $73$. We conclude that, in principle, we can explore the second-order vector mode by using 21cm radiation from the dark ages. By comparing PGWs, it was also found that the PGWs with a tensor-to-scalar ratio $r \approx 10^{-5}$ become subdominant in the 21cm lensing curl mode. In the previous study \cite{Book:2011dz}, they concluded that it is possible to detect the PGWs with $r\approx 10^{-9}$. However, when second-order effects are included in their analysis, a tensor-to-scalar ratio smaller than $r\lesssim 10^{-5}$ would be difficult to detect by the 21cm lensing curl mode. We can generalize this discussion for any model, including the vector or tensor modes with model parameters. The second-order vector mode is generated from the first-order scalar mode that has been well determined by current observations. Therefore, the 21cm curl mode from the second-order vector mode always exists. Even if 21cm lensing is induced by other models, an amplitude smaller than the second-order vector mode is difficult to detect with 21cm lensing. Throughout this paper, we assumed the ideal and challenging experiment for 21cm signals. There are some forthcoming observations for 21cm signals after recombination, e.g., from the Square Kilometer Array. Moreover, exploring 21cm radiation must become an active topic in the near future. Before starting these observations, exploring the potentials of 21cm radiation is important and this work gives one of the nontrivial solutions.

1607.03973

1607

1607.06621_arXiv.txt

{The unprecedented photometric quality and time coverage offered by the {\sl Kepler} spacecraft has opened up new opportunities to search for signatures of nonlinear effects that affect oscillation modes in pulsating stars.} {The data accumulated on the pulsating hot B subdwarf KIC\,10139564 are used to explore in detail the stability of its oscillation modes, focusing in particular on evidences of nonlinear behaviors.} {We analyse 38-month of contiguous short-cadence data, concentrating on mode multiplets induced by the star rotation and on frequencies forming linear combinations that show intriguing behaviors during the course of the observations.} {We find clear signatures that point toward nonlinear effects predicted by resonant mode coupling mechanisms. These couplings can induce various mode behaviors for the components of multiplets and for frequencies related by linear relationships. We find that a triplet at 5760\,$\mu$Hz, a quintuplet at 5287\,$\mu$Hz and a ($\ell>2$) multiplet at 5412\,$\mu$Hz, all induced by rotation, show clear frequency and amplitude modulations which are typical of the so-called intermediate regime of a resonance between the components. One triplet at 316\,$\mu$Hz and a doublet at 394\,$\mu$Hz show modulated amplitude and constant frequency which can be associated with a narrow transitory regime of the resonance. Another triplet at 519\,$\mu$Hz appears to be in a frequency lock regime where both frequency and amplitude are constant. Additionally, three linear combination of frequencies near 6076\,$\mu$Hz also show amplitude and frequency modulations, which are likely related to a three-mode direct resonance of the type \threemodes{}.} {The identified frequency and amplitude modulations are the first clear-cut signatures of nonlinear resonant couplings occurring in pulsating hot B subdwarf stars. However, the observed behaviors suggest that the resonances occurring in these stars usually follow more complicated patterns than the simple predictions from current nonlinear theoretical frameworks. These results should therefore motivate further work to develop the theory of nonlinear stellar pulsations, considering that stars like KIC\,10139564 now offer remarkable testbeds to do so.}

Hot B subdwarf (sdB) stars are helium core burning objects that populate the so-called Extreme Horizontal Branch (EHB). They are expected to have a mass around 0.47 $M_{\odot}$ and are characterized by a very thin hydrogen-rich residual envelope containing at most $\sim 0.02$ $M_{\odot}$. For this reason, they remain hot and compact throughout all their helium core burning evolution, with effective temperatures, $T_{\rm eff}$, and surface gravities, $\log g$, ranging from 22\,000 K to 40\,000 K and from 5.2 to 6.2, respectively \citep{he09,fo12}. The presence of pulsations in some sdB stars make them good candidates for probing their interior with the technique of asteroseismology. A first group of nonradial sdB pulsators with periods of a few minutes was theoretically predicted by \citet{ch96} and effectively discovered by \citet{ki97}. These pulsators, now referred to as the V361\,Hya stars, show low-order, low-degree pressure ($p$-)modes that are driven by a $\kappa$-mechanism induced by the partial ionization of iron-group elements occurring in the "Z-bump" region and powered-up by radiative levitation \citep{ch96,ch97}. Long period oscillations of $\sim1-4$ h were later discovered by \citet{gr03}, forming another group of sdB pulsators known as the V1093~Her stars. The latter show mid-order gravity ($g$-)modes driven by the same mechanism \citep{fo03}. Hybrid pulsators that show both $p$- and $g$-mode oscillations simultaneously have also been reported \citep[e.g., ][]{ss06}. Tight seismic constraints have indeed been obtained from the measured frequencies using both types of sdB pulsators, in particular based on high-quality photometric data gathered from space-borned telescopes \citep[e.g., ][]{ch11b,van10}. However, the reason behind the apparent variability of some oscillation modes in sdB stars, already noticed from repeated ground based campaigns \citep[e.g., ][]{ki07}, has remained poorly understood. The temporal variation of oscillation modes in pulsating sdB stars is beyond the scope of the standard {\sl linear} nonradial stellar oscillation theory in which eigenmodes have a stable frequency and amplitude \citep{un89}. These behaviors must be studied within a {\sl nonlinear} framework to interpret the modulations. In particular nonlinear resonant mode coupling effects are expected to affect some oscillation modes, as noted, e.g., in the helium dominated atmosphere white dwarf variable (DBV) star GD~358 \citep{go98}. Different types of resonant coupling have been investigated within the framework of the amplitude equation (AE) formalism since the 1980's, among them the \threemodes{} resonance \citep{dz82,mo85} and the 2:1 resonance in Cepheid stars \citep{bu86}. The AE formalism was then extended to nonadiabatic nonradial pulsations in Eulerian and Lagrangian formulations by \citet{go94} and \citet{van94}, respectively. A theoretical exploration of specific cases of nonradial resonances was developed in \citet{bu95,bu97}, including notably the resonance occurring in a mode triplet that is caused by slow stellar rotation and which satisfies the relationship $\nu_+ + \nu_- \sim 2\nu_0$ , where $\nu_0$ is the frequency of the central $m=0$ component. However, these theoretical developments based on AEs have since considerably slowed down, in part due to the lack of clear observational data to rely on. The launch of instruments for ultra high precision photometry from space has changed the situation, making it now possible to capture amplitude and/or frequency modulations occurring on timescales of months or even years that were difficult to identify from ground-based observatories. It is however from ground based data that \citet{va11} proposed for the first time that resonant couplings within triplets could explain the long-term variations, both in amplitude and frequency, seen in several oscillation modes monitored in the GW\,Virginis pulsator PG\,0122+200, through successive campaigns. The observation of a multitude of pulsating stars, including sdB and white dwarf stars, by the {\sl Kepler} spacecraft has open up new opportunities to indentify and characterize the mechanisms that could modulate the oscillation modes. {\sl Kepler} monitored a 105 deg$^2$ field in the Cygnus-Lyrae region for around four years without interruption, thus obtaining unprecedented high quality photometric data for asteroseismology \citep{gi10}. These uninterrupted data are particularly suited for searching long-term temporal amplitude and frequency modulations. In the context of white dwarf pulsators, for instance, \citet[hereafter Z16]{zo16} found that \dbv{} shows clear signatures of nonlinear effects attributed to resonant mode couplings. In this star, three rotational multiplets show various types of behaviors that can be related to different regimes of the nonlinear resonant mode coupling mechanism. In particular some amplitude and frequency modulation timescales are found to be consistant with theoretical expectations. This finding suggests that the variations of some oscillation modes in sdB stars may also be related to nonlinear resonance effects. It is in this context that we decided to search clues of similar nonlinear phenomena involving mode interactions in pulsating sdB stars. Eighteen sdB pulsators have been monitored with {\sl Kepler} (see \citealt{os14} and references therein). In this paper, we focus on one of them, the star KIC~10139564, which was discovered in quarter\,Q2.1 and then continuously observed from Q5.1 to Q17.2. A preliminary analysis based on one month of short cadence data originally showed that KIC~10139564 is a V361-Hya type (rapid, $p$-mode) sdB pulsator featuring also a low-amplitude $g$-mode oscillation \citep{ka10}. With extended data, \citet{ba12} detected up to 57 periodicities including several multiplets attributed to the rotation of the star. These multiplets are characterized by common frequency spacings, both for the $p$- and $g$-modes, indicating that KIC~10139564 has a rotation period of $25.6\pm1.8$\,d. These authors did not find any radial-velocity variations from their dedicated spectroscopy and derived the atmospheric parameter values $T_{\rm eff}=31~859$ K and $\log g=5.673$ for this star. An interesting finding concerning KIC~10139564 is that two of the identified multiplets may have degrees $\ell$ greater than 2, a possibility further investigated by \citet{bo13}. The detection of several multiplets in this star continuously monitored for more than three years makes it a target of choice for studying eventual nonlinear resonant mode couplings in sdB stars. In this study, we show that several multiplets in KIC\,10139564 have indeed amplitude and frequency modulations suggesting nonlinear resonant mode couplings, which constitutes the first clear-cut case reported for sdB pulsators, so far. In Sect.\,2, we present the thorough analysis of the frequency content of the {\sl Kepler} photometry available on KIC\,10139564, including our analysis of the frequency and amplitude modulations identified in several multiplets and linear combination frequencies. In Sect.\,3, we recall some theoretical background related to nonlinear resonant mode couplings, focusing mainly on two types of resonances. The interpretation of the observed modulations which may relate to nonlinear resonant mode couplings is discussed in Sect.\,4. The summary and conclusion are then given in Sect.\,5. \begin{figure} \includegraphics[width=8.5cm]{lc_a.eps} \includegraphics[width=8.5cm]{lc_b_o.eps} \caption{{\sl Top panel}: Condensed representation of the full {\sl Kepler} light curve of KIC\,10139564 (Amplitude as the residual relative to the mean brightness intensity of the star vs time in Barycentric Julian Date) covering from Q5.1 to Q17.2 ($\sim 1147.5$ days). {\sl Bottom panel}: Close-up view showing 0.8 days of the {\sl Kepler} light curve by slices of 0.08 days. At this scale the oscillations are clearly apparent. \label{lc}} \end{figure} \begin{figure*} \includegraphics[width=17cm]{LS} \caption{Lomb-Scargle Periodogram (LSP; Amplitude in \% of the mean brightness vs frequency in $\mu$Hz) of the {\sl Kepler} light curve for KIC\,10139564. The represented range, up to the Nyquist frequency, covers the long-period {\sl g}-mode and the short-period {\sl p}-mode frequency domains. The region between the two dashed vertical lines at 5200 and 6400\,$\mu$Hz is where peaks have the largest amplitudes. However, weaker peaks outside of this particular region are present and are made visible by scaling up amplitudes by a factor of 20. The dashed horizontal line represents the 5.6$\sigma$ detection threshold (see text). Some well-known {\sl Kepler} instrumental artefacts are present, but can easily be recognized. \label{lsp}} \end{figure*}

While studying the high-quality and long-duration photometric data provided by the {\sl Kepler} spacecraft on the pulsating sdB star KIC\,10139564, we have identified different patterns in the frequency and amplitude modulations of the oscillation modes belonging to several rotationally split multiplets or linear combination frequencies. These modulations show signatures that can be associated to nonlinear resonant mode coupling mechanisms that could occur between the multiplet components themselves and with other modes under certain conditions, i.e., satisfying a \threemodes{} resonance relationship. This is the first time that such signatures are quite clearly identified in pulsating hot B subdwarf stars, and the second case reported so far for a compact pulsator monitored with {\sl Kepler} photometry (see Z16). We first reanalysed the 38-month of {\sl Kepler} photometry obtained for KIC\,10139564, leading to the detection of 60 independent frequencies above a secured detection threshold (5.6$\sigma$; see Table\,\ref{t2}). Among these, 29 frequencies consist of three triplets, one doublet, one quintuplet and two incomplete multiplets with $\ell>2$ (see Table\,\ref{t1}). Another three detected frequencies are linked to other frequencies through linear combinations. Five additional groups of frequencies are found in the region between 5400 and 6400\,$\mu$Hz, which have very complicated structures. Finaly, we also find 14 independent frequencies and two frequencies satisfying linear combination relationships that could be real as their amplitudes are between 5 and 5.6$\sigma$ above the noise. In general, our well secured frequencies are in good agreement with the former analysis from \citet{ba12}. In this paper, we particularly concentrated our study on six multiplets and three linear combination frequencies observed near 6076\,$\mu$Hz. We found different types of mode behaviors occurring in the above mentioned frequencies. A "short timescale" quasi-periodic amplitude and frequency modulations along with a slow trend of the frequencies to convergence toward each other occur in the dominant $p$-mode triplet near 5760\,$\mu$Hz ($T_1$). The $\sim570$-day quasi-periodic frequency modulation evolve in antiphase between the two side components in this triplet. Modulated frequencies and amplitudes are also found in a quintuplet near 5287\,$\mu$Hz ($Q_1$) and a ($\ell>2$) multiplet near 5412\,$\mu$Hz ($M_1$), but the modulations do not show a clear periodicity. One triplet near 316\,$\mu$Hz ($T_2$) has a distinct behavior from the above mentioned multiplets, as it shows stable frequencies but varying amplitudes. A similar phenomenon occurs in a doublet near 394\,$\mu$Hz ($D_1$) which shows constant frequencies and an $\sim 1100$ days periodic amplitude modulations. Another triplet at 518\,$\mu$Hz ($T_3$) completely differs from all the above multiplets, with constant amplitudes and frequencies throughout the whole observation run. In addition, we also discovered amplitude and frequency variations in three frequencies near 6076\,$\mu$Hz ($C_1$) that are linked to other independent frequencies through linear combinations. After ruling out various possible causes for the modulations, we showed that these mode behaviors could be related to the different types of nonlinear resonances that should occur according to the amplitude equation formalism. In particular, nonlinear resonant couplings within a multiplet can lead to three main regimes, all of which are possibly occurring in KIC\,10139564. The multiplets $T_1$, $Q_1$ and $M_1$ can be associated with the intermediate regime of the resonance where the involved modes have modulated amplitudes and frequencies. The triplet $T_2$ and doublet $D_1$ have a different behavior that could be associated to a narrow transitory regime in which the frequencies of the modes can be locked (constant) while the amplitudes experience modulations. The behavior of the triplet $T_3$ is the unique case found in this star that can be associated to the frequency lock regime of the resonance, where both amplitudes and frequencies are stable. In addition, the large amplitude ratios between the $C_1$ frequencies and their main parent modes, together with the large variation of amplitude and frequency observed for these peaks, suggest that $C_1$ correspond to three-mode direct resonances. We indeed found that the frequencies of $C_1$ exactly follow the evolution of their main parent modes. Moreover, as the parent modes of $C_1$ are also the components of $T_1$ and $T_2$, we suggest part of the complexity of the mode behaviors could be related to these cross interactions between the various modes. In particular, the slow variations occurring in $T_1$ may be related to the \threemodes{} resonance superimposed to the triplet resonance occurring between the components. We emphasize that the observed frequency modulations likely induced by nonlinear mode interactions could challenge future attempts to measure the evolutionary effects on the oscillation mode periods in pulsating sdB stars. Compared to the resonant variations taking place on timescales of years, the rate of period change of the pulsations due to stellar evolution in sdB stars is much longer, typically occurring on a timescale of $\sim 10^6$\,yr \citep{ch02}. Nonlinear modulations of the frequencies can potentially jeopardize any attempt to measure reliably such rates, unless they can be corrected beforehand. These nonlinear modulations could also complicate the detection of exoplanets or stellar companions around sdB stars using the technique of measuring phase changes in the pulsations \citep{si07}. It should be possible however to distinguish between the two effects, considering that nonlinear couplings may induce different behaviors on different modes, while external causes such as an orbiting body should affect all modes similarly. Finally, we note that our analysis suggests that resonances occurring in real stars, in which modes could be involved in two or more types of different couplings, lead to more complicated patterns than those predicted by current theoretical frameworks which treat the modes only as isolated systems within one type of resonance and ignore the nonlinear interactions that could occur simultaneously outside of the system. This should motivate further theoretical work to develop nonlinear stellar pulsation theory for more precise predictions of the mode behaviors in pulsating stars in general.

1607.06621

1607

1607.04412_arXiv.txt

With the aim of paving the road for future accurate astrometry with MICADO at the European-ELT, we performed an astrometric study using two different but complementary approaches to investigate two critical components that contribute to the total astrometric accuracy. First, we tested the predicted improvement in the astrometric measurements with the use of an atmospheric dispersion corrector (ADC) by simulating realistic images of a crowded Galactic globular cluster. We found that the positional measurement accuracy should be improved by up to $\sim2$ mas with the ADC, making this component fundamental for high-precision astrometry. Second, we analysed observations of a globular cluster taken with the only currently available Multi-Conjugate Adaptive Optics assisted camera, GeMS/GSAOI at Gemini South. Making use of previously measured proper motions of stars in the field of view, we were able to model the distortions affecting the stellar positions. We found that they can be as large as $\sim 200$ mas, and that our best model corrects them to an accuracy of $\sim1$ mas. We conclude that future astrometric studies with MICADO requires both an ADC and an accurate modelling of distortions to the field of view, either through an a-priori calibration or an a-posteriori correction.

\label{sec:intro} % Accurate astrometry is one of the major drivers for diffraction limited Extremely Large Telescopes (ELTs). To reach diffraction limited observations, the Multi-AO Imaging Camera for Deep Observations (MICADO), one of the first light instruments for the European-ELT, will be assisted by an Adaptive Optics module (MAORY, \cite{diolaiti10}) providing both a Single Conjugate (developed jointly with the MICADO consortium) and a Multi Conjugate modes. The goal of MICADO is a relative astrometric accuracy for bright and isolated stars of 50 $\mu$as over a central, circular field of 20 arcsec diamater. To determine if such an ambitious goal is feasible, a dual approach is taken. To simulate stellar fields as they would be seen by the SCAO with the predicted instrumental Point Spread Function (PSF) and analyse the resulting realistic images will test the predicted performance, and help to optimise the instrumental design. In addition, present-day astrometric studies with existing MCAO facilities are crucial to test the main sources of inaccuracy not related to the specific instrumental design or telescope. In this paper we present both these approaches, as complementary studies. In particular, we start by investigating in Section \ref{sim} how to best reach the astrometric requirement for MICADO by quantifying the errors associated to one of the most important components in the light path: the atmospheric distortion corrector (ADC). This investigation is carried out making simulations with the SCAO module PSF of the central region of a crowded globular cluster, for a field of view of $2 \times 2$ arcsec. This size is small enough for PSF variations not to be important, but big enough to contain a sufficient number of stars. Then, in Section \ref{real}, we also present the results of an astrometric study performed with Gemini Multi-Conjugate Adaptive Optics System (GeMS) observations of the Galactic globular cluster NGC6681. This cluster is the most centrally concentrated in the Galaxy, and thus represents a major observational challenge in terms of stellar crowding. We have previously determined proper motions by comparing two Hubble Space Telescope epochs (\cite{massari13}). This makes this cluster an ideal candidate to test the effects that will be introduced by MCAO corrections on proper motion measurements and related uncertainties. In particular, we looked for systematic distortions introduced in the GeMS images by observing through both J and Ks filters and we quantify their impact on the astrometry. Though previous studies have already tested the astrometric performance of MCAO cameras (e.g. \cite{neichel14a, ammons14, lu14, fritz16}), our investigation is the first to address the detailed structure and sizes of distortions in a MCAO instrument and will therefore be the starting point for understanding MCAO astrometry capabilities and any future improvements in the calibration and data reduction strategy for ELT observations.

1607.04412

1607

1607.07377_arXiv.txt

{It has been observed \citep{Wilthvar,Wilthvar2} that the locally measured Hubble parameter converges quickest to the background value and the dipole structure of the velocity field is smallest in the reference frame of the Local Group of galaxies. We study the statistical properties of Lorentz boosts with respect to the Cosmic Microwave Background frame which make the Hubble flow look most uniform around a particular observer. We use a very large N-Body simulation to extract the dependence of the boost velocities on the local environment such as underdensities, overdensities, and bulk flows. We find that the observation \citep{Wilthvar,Wilthvar2} is not unexpected if we are located in an underdensity, which is indeed the case for our position in the universe. The amplitude of the measured boost velocity for our location is consistent with the expectation in the standard cosmology. % } \author[a]{David Kraljic} \author[a,b]{\& Subir Sarkar} \affiliation[a]{Rudolf Peierls Centre for Theoretical Physics, University of Oxford,\\ 1 Keble Road, Oxford, OX1 3NP, United Kingdom} \affiliation[b]{Niels Bohr International Academy, University of Copenhagen,\\ Blegdamsvej 17, 2100 Copenhagen, Denmark} \emailAdd{David.Kraljic@physics.ox.ac.uk}\emailAdd{Subir.Sarkar@physics.ox.ac.uk}

The standard Lambda Cold Dark Matter ($\Lambda$CDM) cosmological model is based on an assumed background Friedmann-Lema\^itre-Robertson-Walker (FLRW) geometry that is isotropic and homogeneous. The structure in the universe is modelled as statistically isotropic and homogeneous, initially Gaussian distributed perturbations to this background. This would result in some scatter in the Hubble diagram due to local `peculiar' (non-Hubble) velocities. The cosmic rest frame is defined by comoving observers in the FLRW background and the Cosmic Microwave Background (CMB) frame is conventionally taken to correspond to this `standard of rest' in which both the leading order linear Hubble law and peculiar velocities are defined. The observed dipole pattern in the CMB is then interpreted as due to our peculiar velocity with respect to the cosmic rest frame. For an ensemble of observers, the frame where the velocity flow around an average observer is most uniform corresponds to the CMB frame. However, for a particular observer, the most uniform flow may well be in a frame boosted with respect to the CMB. In this paper we are primarily concerned with the reference frame wherein the Hubble parameter in successive radial shells converges most quickly to the background value. We call this the `minimum Hubble variation frame' or frame with most uniform Hubble flow. % The authors of \citep{Wilthvar,Wilthvar2} looked for a different standard of rest based on the uniformity of the Hubble flow. They found that the Hubble parameter averaged in radial shells converges quickest to its background value in a frame that is boosted with respect to the CMB frame and corresponds roughly to the frame of the Local Group (LG) of galaxies. Moreover, the dipole structure of the velocity flow persists in the CMB frame at large distances, contrary to expectation, but is smaller after boosting to the LG frame. Both of these observations make the local universe as seen in the boosted LG frame closer to the usual ($\Lambda$CDM) expectation. It was suggested \citep{Wilthvar2} that the velocity of the minimum Hubble variation frame should correspond roughly to the group velocity of the `finite infinity' region \citep{EllisFI,WiltFI, WiltFI2}. Ref.\citep{Bolejkodiff} continued the study of the Hubble flow anisotropy with emphasis on the non-kinematic differential expansion of space as the origin of (at least a part) of the CMB dipole. Note that this general relativistic effect where the Hubble parameter is both a function of space and time is \emph{not} captured in N-Body simulations where by construction there is a single background expansion rate. These studies \citep{Wilthvar,Wilthvar2,Bolejkodiff} were in part motivated by some analyses of the local bulk flow of galaxies (as measured in the CMB frame) which show a lack of convergence to the CMB frame even beyond $\sim100$ Mpc \citep{lauerbulk,hudbulk2,kashlinskybulk2,watkinsbulk,lavabulk,kashlinskybulk,feldmanbulk,colinbulk,6dfbulk,2012arXiv1210.0625G,macbulk,watkinsbulk2}, as would be the na\"ive expectation if the universe is indeed homogeneous on larger scales as is inferred from galaxy counts in the SDSS \cite{2005ApJ...624...54H} and WiggleZ \cite{2012MNRAS.425..116S} surveys. Other authors have used different methods (and data) to argue however that observed bulk flows are consistent with $\Lambda$CDM \citep{nusserbulk,daibulk,turnbullbulk,mabulk,mabulk2,hongbulk,applbulk,hutebulk,scrimbulk,feinbulk}. The situation is rather confusing e.g. Ref.\citep{colinbulk} showed using 165 SNe~Ia with redshift $z \lesssim 0.1$ in the Union~2 catalogue that there is an anomalously high and apparently constant bulk flow of $\sim250$~km/s extending out to the Shapely supercluster at $z \simeq 0.06$ ($\sim260$~Mpc). This was confirmed using 117 new SNe~Ia from the Nearby Supernova Factory survey and it was shown that the flow extends out beyond Shapely, nevertheless these authors concluded that their finding is in agreement with $\Lambda$CDM \citep{feinbulk}. SNe~Ia catalogues do have rather incomplete distributions on the sky which can bias the result, however the discrepancy with the standard expectation has been \emph{confirmed} by analysis of the 6dF galaxy redshift survey which is the largest, most homogeneously derived peculiar velocity sample to date \citep{2012arXiv1210.0625G}. The variance of the Hubble parameter can be calculated in linear perturbation theory (e.g. Ref.\citep{wanglin}) or using N-Body simulations (e.g. Ref.\citep{HubbleNBody}). In the latter study the dependence of the local Hubble parameter on the observer's environment (i.e. halos versus voids) was explored and it was found, as expected, that in voids the local value of the Hubble parameter is biased towards higher values while the opposite is true for observers in halos. Given the evidence that we are located in a large under-dense region \citep{hbubble,hbubble2,underdensity,2014MNRAS.437.2146W} this is particularly relevant for explaining the tension between the locally measured value of $H_0$ and the (smaller) value inferred from fitting CMB data (see e.g. Refs.\citep{marrahubble,bendayanhubble}). Another approach is to reconstruct the local universe in N-Body simulations e.g. Ref.\citep{hesshubble} finds that our particular position does bias the locally measured Hubble parameter upwards by about 2\%. However others conclude that the expected variance is far too small to explain the current discrepancy between local and global probes of $H_0$ \citep{2014JCAP...10..028O,2016JCAP...02..001O}. In Ref.\citep{Liuhubble} it is argued that the Hubble parameter can be measured more precisely if observations are restricted to only the zones around critical points of the velocity field, in contrast with the usual approach of increasing statistics by averaging indiscriminately over large datasets. In this paper we study the statistical properties of boosted frames in which the spherically averaged Hubble flow looks most uniform. We are specially interested whether such a boost particular to our position (as measured in Ref.\citep{Wilthvar2}) is consistent with the $\Lambda$CDM expectation. We also look at the dipolar structure of the Hubble flow and its role in determining the frames of minimum Hubble variation. We re-derive the expression for the systematic offset of Hubble parameters between different reference frames paying attention to the dipolar structure. Our expression agrees with the measured offset between the CMB and LG frames and resolves much of the discrepancy found in previous studies \citep{Wilthvar,Wilthvar2}. We use one of the biggest N-Body simulations of $\Lambda$CDM to date \citep{darksky} having a volume of (8\,$h^{-1}$Gpc)$^3$ with the observers placed randomly in large-scale under- and over-densities. We find that the expectations derived from the simulation are consistent with the measurements for our particular position \citep{Wilthvar,Wilthvar2}. Specifically, for an observer in an under-density, a boost simultaneously makes the spherically averaged Hubble parameter converge quicker to the background value and reduces the dipole structure. We also find that the boost velocity of the frame that minimises the Hubble variation for our location is \emph{consistent} with the distribution extracted from the $\Lambda$CDM simulation and is indeed correlated with the group velocity of the `finite infinity' region as has been suggested earlier \citep{Wilthvar2}. In \S~\ref{sec:sims} we describe the N-Body simulation and the halo finder used in this paper. \S~\ref{sec:frames} summarises the methods and the theoretical concepts used in this paper. Our results are presented and discussed in \S~\ref{sec:results}, and we conclude in \S~\ref{sec:conclusion}.

\label{sec:conclusion} We have studied the properties of boost velocities of observers with respect to the CMB frame that make the spherically averaged Hubble flow converge the quickest to its background value. We place observers at different locations in a Hubble volume N-body simulation and find that the distribution of boost velocities they observe is near Gaussian with the amplitude dependent on the location --- the larger the overdensity the larger the amplitude of the typical boost required. For observers in underdense regions, on average, the boosts that make the spherically averaged Hubble parameter converge fastest to the background value reduce at the same time the dipole structure of the Hubble flow. Based only on such local measurements of the velocity field the observers would choose such a frame as the cosmic rest frame given that it is closest to the na\"ive FLRW expectation. However for observers in overdense regions, on average, the boosts that make the Hubble parameter closer to the background value increase the dipole of the velocity field. We show that the boost velocity to the frame of most uniform flow is correlated with the bulk flow velocities and in particular with the group velocity of the `finite infinity' region, as was suggested in \citep{Wilthvar2}. The amplitude of the boost for our position \citep{Wilthvar2} in the universe is found \emph{not} to be in tension with the $\Lambda$CDM expectations. Note that the effects studied here are most pronounced in the mildly non-linear and non-linear regime, i.e. below the $z = 0.023$ threshold adopted \citep{riesshubble} in the determination of $H_0$ from SNe~Ia data. Additionally we re-derived the expression for the systematic offset of the Hubble parameter between different frames, noting that the dipole structure in spherical shells cannot be boosted away entirely. Our expression agrees with the measured difference between $H_\text{CMB}$ and $H_\text{LG}$ and resolves discrepancies in previous work.

1607.07377

1607

1607.03554_arXiv.txt

Stellar streams have become central to studies of the interaction histories of nearby galaxies. To characterize the most prominent parts of the stellar stream around the well-known nearby ($d$ = 17 Mpc) edge-on disk galaxy NGC 5907, we have obtained and analyzed new, deep $gri$ Subaru/Suprime-Cam and 3.6\,$\mu$m {\it Spitzer}/Infrared Array Camera (IRAC) observations. Combining the near-infrared 3.6 $\mu$m data with visible-light images allows us to use a long wavelength baseline to estimate the metallicity and age of the stellar population along a $\sim$~60 kpc long segment of the stream. We have fitted the stellar spectral energy distribution (SED) with a single-burst stellar population synthesis model and we use it to distinguish between the proposed satellite accretion and minor/major merger formation models of the stellar stream around this galaxy. We conclude that a massive minor merger (stellar mass ratio of at least 1:8) can best account for the metallicity of $-0.3$ inferred along the brightest parts of the stream.

\label{s:intro} Major mergers, mergers that take place between galaxies of comparable mass (say stellar or total mass ratios of 1:3 -- 1:5 or larger, with the ratio being the mass of the lower mass galaxy over the mass of the more massive galaxy), have attracted attention in the past \citep*[e.g.,][]{arp66,schweizer82,hibbard96,bush08} because they have features that are easily detected even in shallow imaging observations. However, major mergers may be red herrings when studying the evolution of typical galaxies, many of which may not have undergone a major merger event since $z$ = 1 \citep[e.g.,][]{xu12}. Much more common \citep[e.g.,][]{rodri15} are ``satellite accretion events'' \citep*[stellar or total mass ratio of mostly $\lesssim$~1:50; e.g.,][]{deason16} that are mergers of a dwarf galaxy with a massive parent galaxy. Evidence for such cannibalism has been found around our own Milky Way galaxy, including the Sagittarius Dwarf stream \citep[e.g.,][]{ibata01b,majewski03}, Monoceros Ring \citep{newberg02}, Anticenter Stream \citep{grillmair06b}, Orphan Stream \citep{grillmair06a,belokurov07}, and Styx Stream \citep{grillmair09}. Such streams have also been seen around the Andromeda galaxy \citep[e.g.,][]{ibata01a}. The effects of mergers intermediate in strength between major mergers and satellite accretion events, a.k.a. ``minor mergers'' (stellar or total mass ratios of roughly between 1:50 and 1:5), have proven to be much harder to observe directly. Such mergers are not expected to transform galaxies as drastically as major mergers (disk to elliptical transformation), but on the other hand one would expect the effects of such mergers to show up as more than just tenuous stellar population changes in the halo (cf. Local Group streams mentioned above). Moreover, \citet{stewart09} have argued that every large galaxy has undergone at least one minor merger during its lifetime. Thus, it can be argued that it is the minor mergers that are the most dynamically significant merging events, in view of their relatively high frequency (see also \citeauthor{zaritsky97} \citeyear{zaritsky97}, \citeauthor{bullock05} \citeyear{bullock05}) and the substantial dynamical and structural impacts on the primary galaxy that such interactions can bring about. These impacts include the heating and thickening of host galaxy disks, growth of galactic bulges, hierarchical build-up of galaxy mass, counterrotating cores in host galaxies, and triggering and maintenance of bar and spiral structures \citep*[e.g.,][]{mori08,purcell10}. Observational evidence for minor mergers has been surfacing during the last decade, as fairly high surface brightness stellar streams around several nearby galaxies have been detected in ground-based visible-light observations \citep[e.g.,][]{shang98,delgado08,delgado09,delgado10,duc15,jennings15}. \begin{figure*} \centering \includegraphics[angle=270,scale=0.7]{f1.pdf} \caption{Observed (left) and modeled (right) stellar stream around NGC 5907, reproduced from Figures 2 and 4 of ``The Ghost of a Dwarf Galaxy: Fossils of the Hierarchical Formation of the Nearby Spiral Galaxy NGC 5907'' by D. Mart\'{i}nez--Delgado, J. Pe\~{n}arrubia, R. J. GaBany, I. Trujillo, S. R. Majewski, \& M. Pohlen, ApJ, vol 689, issue 1, (2008) pp. 184--193. The image on the left is a combined clear luminance (350--850 nm), red, blue, and green filter image, with a color image made from the red, blue, and green filter images superimposed on the saturated disk (see \citeauthor{delgado08} \citeyear{delgado08} for more information), observed at the BlackBird Remote Observatory's 0.5-m telescope. The image on the right is an N-body minor merger model from \citet{delgado08}. The different colors designate particles that become unbound after different pericenter passages, with cyan, green, red, and blue indicating particles that became unbound during the first, second, third, and fourth most recent pericenter passages. The various letter designations are discussed in \citet{delgado08}. q = 1.0 refers to a spherical halo model. The image on the left is about $17\farcm 5$ $\times$ $27\farcm 5$ in size. North is to the right and east is up.\label{visfig}} \end{figure*} NGC~5907 is a nearby ($d$ = 17 Mpc; \citeauthor{tully13} \citeyear{tully13}, AB$_{\rm mag}$$^{3.6}$ = $-21.9$; \citeauthor{irsa}, $V_{\rm rot}^{\rm max}$ = 240 km~sec$^{-1}$; \citeauthor{casertano83} \citeyear{casertano83}) edge-on Sc-type disk galaxy with a stellar stream around it, discovered by \citet{shang98}, and studied further by \citet{zheng99} and \citet{delgado08}. The outer disk of NGC~5907 also has a warp \citep[e.g.,][]{shang98}. \citet{delgado08} obtained very deep images of NGC~5907 and revealed the spectacular long stellar stream (Figure~\ref{visfig}). However, because they used a luminance filter \citep[see Fig. 1 in][]{delgado15}, \citet{delgado08} were not able to measure colors in the stream. \citet{delgado08} also produced a simple N-body model that mimicked the observed loopy stellar stream as a fossil of the tidal disruption of a single satellite in a merger event (total mass ratio of 1:4000), rejecting the hypothesis of multiple merger events in the halo of NGC 5907. In contrast, \citet{wang12} reported a model that reproduced the general structure of the stellar stream around NGC~5907 in a major merger scenario. \citet{wang12} suggested that the colors of the stream may be used to distinguish between a minor satellite accretion event and a major merger origin of the stellar stream in NGC~5907 (the stream is too faint to be observed spectroscopically). They used both 1) the color-inferred iron abundance in the disk outskirts (affected also by stars in the inner halo) and found that the $R-I$ color from \citet{zheng99} is similar to that of massive disk galaxies instead of dwarf satellite galaxies, and 2) the mass--metallicity relation in which, for example, the low-mass Sagittarius satellite galaxy (with low [Fe/H]) would produce a bluer color than what is observed. We exploit a similar technique in this paper. We have obtained new deep Subaru/Suprime-Cam $gri$ and {\it Spitzer}/Infrared Array Camera (IRAC) 3.6\,$\mu$m images of NGC 5907 to measure the spectral energy distribution (SED) and color indices of the brightest parts of the stellar stream east of the disk of NGC 5907, and to compare them to the predictions of satellite galaxy accretion and minor/major merger models. To our knowledge this is the first time that IRAC data have been used to constrain the metallicity and the age of the stellar population in a stellar stream around a nearby galaxy (outside the Local Group). This paper is organized as follows. In Section~\ref{obs} we summarize the new observations and data, in Section~\ref{magcolor} we present the surface brightness and color index results and study the potential effects of an extended point spread function (PSF), while Section~\ref{sed} addresses how we fit single burst SED models from the flexible stellar population synthesis (FSPS) code. We discuss our results in Section~\ref{discus} and present our conclusions in Section~\ref{concl}.

\label{concl} We have studied the brightest part of the stellar stream around NGC~5907 with both visible-light ($g$, $i$, and $r$) and IRAC near-infrared 3.6 $\mu$m images. While the stream is very faint, we have managed to generate color index profiles and an SED for the brightest part of the stream. Our results show a slight color change in some of the bands along the stream, most notably in $g-i$ and $r-i$ in which the color becomes bluer after the turning point in the northeast corner of the brightest parts of the stream. However, there may be a remaining background subtraction problem in the $i$-band and our results are essentially consistent with no color gradients along the stream. We have fitted FSPS stellar population synthesis models with a single-age stellar population. The best fit is achieved for an old and relatively metal-rich stellar population in the stream. Such a stellar population appears to be more consistent with a rather massive minor merger event than with the stripped stars from a dwarf galaxy accretion event. Future work should include better visible light photometry of the stream, and improving the stellar population synthesis model predictions in the near-infrared. Future kinematical studies of globular clusters associated with this stream could shed light on the implied minor merger scenario (Alabi et al., in preparation).

1607.03554

1607

1607.03886_arXiv.txt

We present a comprehensive study of the active dM4e star GJ 1243. We use previous observations and ground-based echelle spectroscopy to determine that GJ 1243 is a member of the Argus association of field stars, suggesting it is $\sim 30-50$ Myr old. We analyze eleven months of 1-minute cadence data from \textit{Kepler}, presenting \textit{Kepler} flare frequency distributions, as well as determining correlations between flare energy, amplitude, duration, and decay time. We find that the exponent $\alpha$ of the power-law flare energy distribution varies in time, primarily due to completeness of sample and the low frequency of high-energy flares. We also find a deviation from a single power law at high energy. We use ground-based spectroscopic observations simultaneous with the \textit{Kepler} data to provide simultaneous photometric and spectroscopic analysis of three low-energy flares, the lowest-energy dMe flares with detailed spectral analysis to date on any star. The spectroscopic data from these flares extend constraints for radiative hydrodynamic (RHD) flare models to a lower energy regime than has previously been studied. We use this simultaneous spectroscopy and \textit{Kepler} photometry to develop approximate conversions from the \textit{Kepler} bandpass to the traditional $U$ and $B$ bands. This conversion will be a critical factor in comparing any \textit{Kepler} flare analyses to the canon of previous ground-based flare studies.

M dwarfs are well known for their high magnetic activity, most notably their powerful, frequent flares. These are believed to be magnetic reconnection events analogous to the flares observed on the Sun, but occur much more frequently, and with much greater energies. Observational studies of flares on these stars have been limited by the difficulty of collecting statistically complete samples of detailed flare light curves from ground-based observation. To mitigate the stochastic occurence rate of flares, studies have often chosen to follow multiple active stars with similar spectral types over several nights \citep[e.g.][]{Moffett1974, Hilton2011}. Automated surveys provide repeated imaging of the sky and can efficiently yield millions of individual measurements of M dwarfs, but they do not provide temporal information for individual flares, or observe a complete sample of flares for individual stars. Dedicated ground- and space-based exoplanet surveys, however, provide sufficiently high-cadence observations over a duration long enough to ensure a statistically complete sample of flares on observed stars. Due to the inclusion of local M dwarfs in transiting exoplanet surveys such as MEarth \citep{Nutzman2008} and \textit{Kepler} \citep{Borucki2010}, detailed study of M dwarf activity in the solar neighborhood has dramatically increased in recent years. The \textit{Kepler} survey launched a new era of stellar photometric investigation, allowing for unprecedented light curve collection. While the mission's primary goal has been the detection of exoplanets, its near-continuous short- and long-cadence monitoring of targets make it nearly ideal for recording statistically complete samples of stellar variability, including flares, throughout the main sequence \citep{Basri2010, Walkowicz2011}. Previous works have utilized \textit{Kepler}'s unique capabilities to explore M dwarf activity with unparalled depth and completeness. \citet[hereafter Paper 1]{Hawley2014} analyzed the flare frequency distribution of three active and three inactive M dwarfs with two months of high-cadence \textit{Kepler} data, finding strong correlations between flare energy, amplitude, duration, and decay time, and a weak correlation with rise time. No correlation was found between flare energy or occurrence and starspot phase, and energies of consecutive flares. \citet[hereafter Paper 2]{Davenport2014} expanded on the sample from Paper 1, producing a 90\% complete sample of over 6100 flares from eleven months of \textit{Kepler} short-cadence data collected on GJ 1243. \citet[hereafter Paper 3]{Lurie2015} used \textit{Kepler} data on the active stars GJ 1245 A and B to study the evolution of flares and starspots in a binary system, providing a detailed photometric analysis of a multiple system of fully convective M dwarfs and yielding an important constraint on stellar age-rotation-activity models. Paper 3 also found that some flares for GJ 1245 were bright enough that their peak fluxes equaled or exceeded 95\% of \textit{Kepler}'s quoted full well depth, potentially entering the non-linear regime and saturating the CCD. This paper, the fourth in this series, focuses on the properties of GJ 1243. GJ 1243 is a dMe flare star located in the solar neighborhood \citep[see e.g.][]{Reid2004}. \citet{Hawley1995} and \citet{Hawley1996} identified the star as a dMe based on strong H$\alpha$ emission (EW $>$ 1 \AA) as part of a survey of the preliminary version of the Third Catalog of Nearby Stars \citep{GJ1991}, while \citet{Gershberg1999} later identified the star as a a UV Ceti-type flare star, and \citet{Gizis2002} confirmed it as active. The MEarth survey found a distance to the star of $13.48 \pm 0.42$ pc based on trigonometric parallax \citep{Dit2014}. Additionally, \citet{LepineShara2005} found that GJ 1243 was a high-proper motion star, with a root mean square proper motion of 326 $\mathrm{mas\,yr^{-1}}$. More recently, \citet{Irwin2011} found a period for GJ 1243 of 0.593 days using MEarth photometry. \citet{Savanov2011} used long-cadence data from \textit{Kepler} Q1 to confirm a period of 0.5926 days, and identified two independent regions of starspots that persisted in the same location throughout the observation. \citet{Reinhold2013} found no evidence of differential rotation in GJ 1243 as part of a larger survey of \textit{Kepler} targets, but \citet{Davenport2015} found evidence of differential rotation using phase-tracking and spot modeling. \citet{Ramsay2013} also noted frequent flares on GJ 1243 as a comparison to KIC 5474065 (also an M4V star), and conducted a flare rate analysis based on short-cadence data from \textit{Kepler} Q6 (also used in Paper 1). In this paper, we combine the techniques of all previous work in the \textit{Kepler} Flares series to provide a comprehensive analysis of the behavior of GJ 1243 for the full eleven months of short cadence data collected. In Section 2, we review the methods of observing and recording both photometric and spectroscopic data, and discuss how flares are identified in both data sets. In Section 3, we discuss the kinematics and intrinsic characteristics of GJ 1243: its rotational velocity, proper motion in the galaxy, magnetic field, and metallicity. In Section 4, we analyze the full flare sample from \textit{Kepler}, investigating the flare frequency-energy relationship and the correlations found in Paper 1. We also examine these in shorter (month-by-month and quarter-by-quarter) time units, searching for predictable changes over time. In Section 5, we examine spectroscopic observations of GJ 1243 collected simultaneously with \textit{Kepler} observations to break modeling degeneracies that result from the single bandpass on \textit{Kepler}, and characterize the spectral behavior of the observed flares. Finally, in Section 6, we discuss these observations and build a full characterization of GJ 1243.

\label{sec:Conclusions} In this paper, we have presented a comprehensive study of the star GJ 1243, using previous data in the literature, 11 months of short-cadence wide-band photometry from \textit{Kepler}, and simultaneous spectroscopy for three flares. We calculated and analyzed the cumulative frequency of flares as a function of energy; the relationships between flare energy, amplitude, and duration; the relationship of flare frequency and energy with stellar phase; and several characteristics of flares based on spectroscopy. In the following sections, we discuss the potential for a secondary power law for frequency of the highest-energy flares, potential origins of the flare morphology observed in the \textit{Kepler} data based on the spectral data, and potential extensions of the discoveries made here to other stars. \subsection{Basic Characteristics of the Star} We demonstrated based on high-resolution echelle spectroscopy that GJ 1243 has a large $v \sin i$ of 25 $\mathrm{km s^{-1}}$, as expected for its high activity level. We used the observed period in its \textit{Kepler} light-curve and this measurement to find that the system's inclination angle $i$ was $32^{\circ}$. Using data on its proper motion, distance, and radial velocity, we showed that GJ 1243 is likely (98.5\% probability) a member of the Argus association, suggesting that its age is 30-50 Myr. This relatively young age agrees well with the high level of activity we observe on this star. \subsection{Frequency of High-Energy Flares} We find that generally a single power-law model fits each data set in monthly and quarterly increments, as well as the full data set. Over time, we show that $\alpha$ shows some uncorrelated variability over time, likely due to the varying sample of flares over each epoch studied. Interestingly, the values of $\alpha$ for each subset are generally lower than the value for the full dataset. We believe this is in part due to the sample size of the full data set allowing us to probe the FFD at a higher energy resolution, compared to the subsets. Additionally, we find that the highest energy flares, both saturated and unsaturated, deviate from a single power law in a way that suggests a steeper power law than the full sample. This would have significant implications for stellar modeling, as values of $\alpha > 2$ are a possible mechanism for the heating of the solar corona. However, these heating models require this larger value of $\alpha$ to extend to low energies, which our data do not. For the full data set $\alpha \approx 2$, though there is a turn-off at low energies where flares above the minimum energy threshold are not detected as expected, as was observed in Paper 1. Paper 2 showed that the 11-month sample used here is 90\% complete at the highest energy levels, suggesting that there may be ``missing'' flares from the sample that could account for this. Additionally, some lower-energy ``classical'' events could have been incorporated into ``complex'' events. While Paper 2 found that this potential overlap could not account for all of the ``missing'' portion, it could be an additional contributing factor. \subsection{Correlations of Flare Characteristics} We find a strong, positive correlation between flare energy and flare duration, as well as between flare energy and amplitude, similar to the trends found in Paper 1. We find that complex flares have higher energies and longer durations than classical flares at the same amplitude. We quantify the relationship between duration and energy, and find that complex and classical flares show a distinct difference in this relationship, which agrees with the characterization of some complex flares as a superposition of classical flares, presented in Paper 2. As with Paper 1, we find no correlation of flare timing or energy with the phase caused by the persistent starspots on the surface, in monthly or quarterly time periods or in the full data set. This suggests that flaring is distributed across the stellar surface, rather than concentrated only in the regions where these starspots appear. Additionally, the high-latitude constraint for the persistent spot found in \citet{Davenport2015}, combined with the inclination angle found in Section 3, indicates that there would be minimal correlation between flare occurrence and phase were the flares in fact associated with the spot. \subsection{Flare Spectroscopy} We observed three low-energy flares with simultaneous ground-based spectroscopy and \textit{Kepler} photometry. Using this simultaneous photometric and spectroscopic data, we were able to classify the third flare of this set as a ``gradual-flare'' type event, following the definition from \citet{Kowalski2013}. We were also able to derive an impulsiveness index $\mathcal{I}$ for this flare based on the \textit{Kepler} data, finding that this value was two orders of magnitude lower than the same index in the $U$ band (synthesized from spectra), due to the smaller flare-induced flux enhancement in the \textit{Kepler} bandpass. We also presented a spectroscopy-based explanation for the two-component empirical flare model adopted in Paper 2, finding that the gradual phase HB and blackbody were tied to each other, but resulted from a cooling region physically distinct from that of the impulsive phase. This suggests that the impulsive and gradual phases are physically distinct phenomena potentially arising from the same initial conditions, which allows for the possibility of separately modeling the impulsive and gradual phases of the flare. \subsection{Application to Future Flare Studies} Using simultaneous photometric and spectroscopic data, we were able to develop a relationship between flux recorded in the \textit{Kepler} bandpass and in the more traditional $U$ and $B$ bandpasses at the peak of a faint flare, directly connecting this result to the larger canon of previous flare surveys. This basic conversion will be invaluable for future studies of dMe flares recorded with \textit{Kepler}, and will be improved by future work that will approach the problem beginning with the empirical \textit{Kepler} flare model.

1607.03886

1607

1607.03909_arXiv.txt

The evolution of the number density of galaxies in the universe, and thus also the total number of galaxies, is a fundamental question with implications for a host of astrophysical problems including galaxy evolution and cosmology. However there has never been a detailed study of this important measurement, nor a clear path to answer it. To address this we use observed galaxy stellar mass functions up to $z\sim8$ to determine how the number densities of galaxies changes as a function of time and mass limit. We show that the increase in the total number density of galaxies ($\phi_{\rm T}$), more massive than M$_{*} = 10^{6}$ \solm, decreases as $\phi_{\rm T} \sim t^{-1}$, where $t$ is the age of the universe. We further show that this evolution turns-over and rather increases with time at higher mass lower limits of M$_{*}>10^{7}$ \solm. By using the M$_{*}=10^{6}$ \solm lower limit we further show that the total number of galaxies in the universe up to $z = 8$ is $2.0^{+0.7}_{-0.6} \times 10^{12}$ (two trillion), almost a factor of ten higher than would be seen in an all sky survey at Hubble Ultra-Deep Field depth. We discuss the implications for these results for galaxy evolution, as well as compare our results with the latest models of galaxy formation. These results also reveal that the cosmic background light in the optical and near-infrared likely arise from these unobserved faint galaxies. We also show how these results solve the question of why the sky at night is dark, otherwise known as Olbers' paradox.

When discovering the universe and its properties we are always interested in knowing absolutes. For example, it is of astronomical interest to calculate how many stars are in our Galaxy, how many planets are surrounding these stars (Fressin et al. 2013), the total mass density of the universe (e.g., Fukugita \& Peebles 2004), amongst other absolutes in the universe's properties. One of these that has only been answered in a rough way is the total number density evolution of galaxies, and thus also the total number of galaxies in the universe. This question is not only of passing interest as a curiosity, but is also connected to many other questions in cosmology and astronomy. The evolution of the number densities of galaxies relates to issues such as galaxy formation/evolution through the number of systems formed, the evolution of the ratio of giant galaxies to dwarf galaxies, the distant supernova and gamma-ray burst rate, the star formation rate of the universe, and how new galaxies are created/destroyed through mergers (e.g., Bridge et al. 2007; Lin et al. 2008; Jogee et al. 2009; Conselice et al. 2011; Bluck et al. 2012; Conselice 2014; Ownsworth et al. 2014). The number of galaxies in the observable universe also divulges information about the mass density of the universe, background light at various wavelengths, as well as insights into Olbers' Paradox. However, there still does not yet exist a good measurement for this fundamental quantity. Understanding the co-moving number density evolution of galaxies has only been possible in any meaningful way since deep imaging with telescopes began with the advent of CCD cameras. Deep surveys to search for distant galaxies started in the 1990s (e.g., Koo \& Kron 1992; Steidel \& Hamilton 1992; Djorgovski et al. 1995), and reached our current depths after deep Hubble Space Telescope imaging campaigns were carried out, especially within the Hubble Deep Field (Williams et al. 1996). This was later expanded to other deep fields such as the Hubble Deep Field South (Williams et al. 2000), the Great Observatories Origins Survey (Giavalisco et al. 2004), and the near-infrared CANDELS fields (Grogin et al. 2011; Koekemoer et al. 2011), and finally to the Hubble Ultra Deep Field (Beckwith et al. 2006) which remains the deepest image in the optical and near-infrared of our universe taken to date. However, despite these surveys it is still uncertain how the total number density of galaxies evolves over time. This is an interesting question as we know that the star formation rate rises, and then declines at $z < 8$ (e.g., Bouwens et al. 2009; Duncan et al. 2014; Madau \& Dickinson 2014), while at the same time galaxies become larger and less peculiar (e.g., Conselice et al. 2004; Papovich et al. 2005; Buitrago et al. 2013; Mortlock et al. 2013; Lee et al. 2013; Conselice 2014; Boada et al. 2015). However, we do not know how the total number of galaxies at a given epoch evolves, and how this is associated with the general formation of the galaxy population as a whole. There are a few reasons for why deep imaging programs are not easily able to convert observations to total numbers of galaxies. One of these issues is that all deep observations are incomplete. This is due to limitations in exposure times and depth such that certain galaxies will be detected more readily than other galaxies. The result of this is an incompleteness down to the magnitude limit of even the deepest surveys, which can be corrected for, but which still leaves some uncertainty. However, the more important issue is that these observations do not reach the faintest galaxies, although from number density fits and theory we know that there should be many more faint galaxies beyond our current observational limits. It is also important to address what we mean by the total number density of galaxies in the universe. This is not a simple quantity to define as the total number density which exists now, the total number density which is observable in principle, and the total number density which is observable with current technology, are all different questions with different answers. There is also the issue that we are limited by the cosmological horizon over what we can observe, and therefore there are galaxies we cannot see beyond it. Even the number of galaxies which exist in the universe today, i.e., if we could view the entire universe as is, and not be limited by light travel time, is a complicated question. Galaxies in the distant universe have evolved beyond what we can currently observe due to the finite nature of the speed of light, and presumably would look similar to those in the local universe. We address these issues in the paper. Our default and ultimate total number density of galaxies we investigate in this paper is how the number density evolves within the current observable universe up to $z \sim 8$. For comparison purposes, we also carry out an analysis in the Appendix of the number of galaxies which are visible to modern telescopes, at all wavelengths, that we can currently observe. We then compare this to measurements of the total number that actually can be potentially observed in the universe based on measured mass functions. We also discuss how these results reveal information concerning galaxy evolution and background light. We also give indications for future surveys and what fraction of galaxies these will observe. This paper is divided up into several sections. \S 2 describes the data we use throughout this analysis, \S 3 describes the results of this paper including using fits of galaxy stellar mass functions to derive the total number of galaxies which are in the universe. \S 4 describes the implications of these results and \S 5 is a summary. Throughout this paper we use a standard cosmology of H$_{0} = 70$ km s$^{-1}$ Mpc$^{-1}$, and $\Omega_{\rm m} = 1 - \Omega_{\lambda}$ = 0.3.

We have investigated the fundamental question of the number density evolution of galaxies in the universe. We research this problem in a number of ways, and discuss the implications for galaxy evolution and cosmology. We use recently measured mass functions for galaxies up to $z \sim 8$ to determine the number density evolution of galaxies in the universe. Our major finding is that the number densities of galaxies decrease with time such that the number density $\phi_{\rm T} (z) \sim t^{-1}$, where $t$ is the age of the universe. We further discuss the implications for this increase in the galaxy number density with look-back time for a host of astrophysical questions. Integrating the number densities, $\phi_{\rm T}$ we calculation that there are $(2.0^{+0.7}_{-0.6}) \times 10^{12}$ galaxies in the universe up to $z = 8$ which in principle could be observed. This is roughly a factor of ten more than is found through direct counting (see Appendix). This implies that we have yet to detect a large population of faint distant galaxies. In terms of astrophysical evolution of galaxies, we show that the increase in the integrated mass functions of all galaxies with redshift can be explained by a merger model. We show that a simple merger model is able to reproduce the decline in the number of galaxies with a merger time-scale of $\tau = 1.29\pm0.35$ Gyr. The derived merger rate at $z = 1.5$ is $R \sim 0.05$ mergers Gyr$^{-1}$ Mpc$^{-3}$, close to the value found through structural and pair analyses. Most of these merging galaxies are lower mass systems based on the increase in number densities with time seen at lower limit selections, using higher masses, for the total number density calculation. We finally discuss the implications of our results for future surveys. We calculate that the number counts of galaxies at magnitudes fainter than $m_{\rm max} = 29$ will largely probe the lower mass galaxies at higher redshifts and eventually at $m_{\rm max} \sim 35$ will turn over and decline due to reaching the limit of the number of galaxies in the universe, unless the mass limit for galaxies is much less than M$_{*} = 10^{6}$ \solm or there are many galaxies at $z > 12$. We also show that this leads to a natural confusion limit in detection and that these galaxies are likely responsible for the optical and near-infrared background and provide a natural explanation for Olbers' paradox. This large additional number of galaxies is also consistent with recent measures of the cosmic infrared background light (e.g., Mitchell-Wynne et al. 2015). In the future, as mass functions become better known with better SED modeling and deeper and wider data with JWST and Euclid/LSST, we will be able to measure the total number densities of galaxies more precisely and thus obtain a better measure of this fundamental quantity. We thank Neil Brandt, Harry Teplitz, and Caitlin Casey for useful discussions concerning non-optical deep detections in galaxy surveys. This work was supported by grants from the Royal Astronomical Society, STFC and the Leverhulme Trust. Support was also provided by NASA/STScI grant HST-GO11082. A.M. acknowledges funding a European Research Council Consolidator Grant. \appendix

1607.03909

1607

1607.03412_arXiv.txt

The \ac{SST-1M} is one of the three proposed designs for the \acp{SST} of the \ac{CTA} project. The \ac{SST-1M} will be equipped with a 4~m-diameter segmented mirror dish and an innovative fully digital camera based on \acp{SiPM}. Since the \ac{SST} sub-array will consist of up to 70 telescopes, the challenge is not only to build a telescope with excellent performance, but also to design it so that its components can be commissioned, assembled and tested by industry. \newline In this paper we review the basic steps that led to the design concepts for the \ac{SST-1M} camera and the ongoing realization of the first prototype, with focus on the innovative solutions adopted for the photodetector plane and the readout and trigger parts of the camera. In addition, we report on results of laboratory measurements on real scale elements that validate the camera design and show that it is capable of matching the \ac{CTA} requirements of operating up to high-moon-light background conditions.

\label{sec:intro} The \ac{CTA}, the next generation very high energy gamma-ray observatory, is a project to build two arrays of over 100 \acp{IACT} placed in two sites in the northern and southern hemispheres. The array will consist of three types of telescopes: \acp{LST}, with $\sim$24~m mirror diameter, \acp{MST}, with $\sim$12~m mirror diameter and small size telescopes (SSTs), with $\sim$4~m mirror diameter\footnote{In this paper, we will call ``mirror'' the full reflective surface of the telescope, which, in our case, is composed of 18 hexagonal facets.}. About 70 small size telescopes will be installed in the southern site, which offers the best view of the galactic plane, and will be spaced at inter-telescope distances between 200-300~m to cover an air shower collecting surface of several km$^2$. This surface allows for observation of gamma-rays with energy between about 3~TeV and 300~TeV~\cite{CTAconcept}. Different \ac{SST} designs are being proposed, among which a single mirror Davies-Cotton telescope (SST-1M) based on \acl{SiPM} (\acs{SiPM}) photodetectors, whose camera is described in this paper. The other two projects~\cite{ASTRI,GCT} are dual mirror telescopes of Schwarzschild-Couder design. \newline \begin{figure*}[htp!] \centering \includegraphics[width=0.7\textwidth]{CTA_CameraAss1.pdf} \caption{CAD drawing indicating the dimensions of the camera.} \label{fig:camera_dim} \end{figure*} The camera is a critical element of the proposed \ac{SST-1M} telescope, and has been designed to address the \ac{CTA} specifications on the sensitivity of the array, its angular resolution, the charge resolution and dynamic range of single cameras, the \ac{FoV} of at least $9^\circ$ for \acp{SST}, the uniformity of the response, as well as on the maintenance time and availability, while keeping the cost of single telescopes reasonable. The SST-1M camera has been designed to achieve the best cost over performance ratio while satisfying the stringent \ac{CTA} requirements. Its components are made with standard industrial techniques, which make them reproducible and suited for large scale production. For these reasons, the camera features a few innovative strategies in both the optical system of the \ac{PDP} and the fully digital readout and trigger system, called DigiCam. \newline A camera prototype is being produced by the University of Geneva - UniGE (in charge of the PDP and its front-end electronics, the cooling system, the mechanics including the shutter, and the system for the integration on the telescope structure), the Jagellonian University and the AGH University of Science and Technology in Krak\'ow (in charge of the development of the readout and trigger system). This prototype not only serves to prove that the overall concept can meet the expected performance, but also serves as a test-bench to validate the production and assembly phases in view of the production of twenty SST-1M telescopes. This paper is structured as follows: the general concept of the camera is described in Sec.~\ref{sec:design}, while Sec.s~\ref{sec:PDP} and \ref{sec:DigiCam} are dedicated to more details on the design of the PDP and of DigiCam, respectively. Sec.~\ref{sec:cooling} describes the cooling system and Sec.~\ref{sec:safety} the housekeeping system. Sec.s~\ref{sec:tests} and \ref{sec:validation} are devoted to the description of the camera tests and validation of its performance estimated with the simulation described in Sec.~\ref{sec:performances}. Sec.~\ref{sec:calibration} describes initial plans on the calibration strategy during operation. In Sec.~\ref{sec:concl}, we draw the conclusions of the results and the plans for future operation and developments.

\label{sec:concl} The prototype camera proposed for the \ac{SST-1M} telescope of the \ac{CTA} project adopts several innovative solutions conceived to provide high performance and reliability on long time scales, as well as being cost effective in view of a possible production of up to 20 units. The challenges encountered during the design phase (such as the realization and operation of large area hexagonal \acp{SiPM} and the hollow light concentrators, the stabilization of the working point of the sensors and the cooling strategy) have been all successfully addressed, and the camera is now being assembled at the University of Geneva, were it will be fully tested and characterized. Preliminary measurements and simulations have shown that the camera fully complies with the CTA requirements. Installation on the prototype telescope structure hosted at the H. Niewodnicza\'nski institute of Nuclear Physics in Krakow is foreseen in fall 2016.

1607.03412

1607

1607.01403_arXiv.txt

Strong gravitational lenses with measured time delays between the multiple images allow a direct measurement of the time-delay distance to the lens, and thus a measure of cosmological parameters, particularly the Hubble constant, $H_{0}$. We present a blind lens model analysis of the quadruply-imaged quasar lens \hequad~using deep {\it Hubble Space Telescope} imaging, updated time-delay measurements from the COSmological MOnitoring of GRAvItational Lenses (COSMOGRAIL), a measurement of the velocity dispersion of the lens galaxy based on Keck data, and a characterization of the mass distribution along the line of sight. \hequad~is the third lens analyzed as a part of the $H_{0}$ Lenses in COSMOGRAIL's Wellspring (H0LiCOW) project. We account for various sources of systematic uncertainty, including the detailed treatment of nearby perturbers, the parameterization of the galaxy light and mass profile, and the regions used for lens modeling. We constrain the effective time-delay distance to be $\tdist = \dt$, a precision of \dtprec\%. From \hequad~alone, we infer a Hubble constant of $H_{0} = \ulcdm$ assuming a flat $\Lambda$CDM cosmology. The cosmographic inference based on the three lenses analyzed by H0LiCOW to date is presented in a companion paper (H0LiCOW Paper V).

\label{sec:intro} The flat $\Lambda$CDM cosmological model is the concordance model of our Universe today. It is consistent with a variety of independent experiments, including an analysis of the cosmic microwave background (CMB) by the {\it Planck} mission \citep{planck2015}. The {\it Planck} results provide the most precise cosmological parameter constraints to date, under the assumption of spatial flatness. However, there is no physical reason to assume flatness, and if the flatness assumption is relaxed, there are strong degeneracies among the cosmological parameters inferred from CMB data, particularly with the Hubble constant, $H_{0}$ \citep[e.g.,][]{freedman2012,riess2016}. Therefore, an independent determination of $H_{0}$ is crucial for understanding the nature of the Universe \citep[e.g.,][]{hu2005,suyu2012c,weinberg2013}. The idea of using gravitational lens time delays to measure the Hubble constant dates back to \citet{refsdal1964}. In practice, gravitational lens time delays provide a one-step method to determine the distance and hence the Hubble constant \citep[e.g.,][]{vanderriest1989,keeton1997,schechter1997,kochanek2003,koopmans2003,saha2006,oguri2007,fadely2010,suyu2010b,suyu2013,sereno2014,rathnakumar2015,birrer2016,chen2016}. This method is independent of the cosmic distance ladder \citep[e.g.,][]{riess2011,freedman2012} and serves as a key test of possible systematic effects in individual $H_{0}$ probes. This method rests on the fact that light rays emitted from the source at the same instant will take different paths through spacetime at each of the image positions. These paths have different lengths and traverse different gravitational potentials before reaching the observer, leading to an offset in arrival times. If the source exhibits variations in its flux, the delays can be measured by monitoring the lensed images. The measured time delays can be used to calculate the time-delay distance, a combination of angular diameter distances among the observer, lens, and source. The time-delay distance is primarily sensitive to $H_{0}$, with weaker dependence on other cosmological parameters \citep[e.g.,][]{coe2009,treu2016}. However, a precise and accurate determination of $H_{0}$ through this method requires a variety of observational data. A dedicated long-term monitoring campaign is necessary to obtain accurate time delays, as the uncertainty in $H_{0}$ is directly related to the relative uncertainty in the measured time delays. Deep, high-resolution imaging is required to accurately model the lens using the extended source images, which is needed to break degeneracies between the mass profile and the underlying cosmology \citep[e.g.,][]{kochanek2002,warren2003}. In order to reduce the effects of the mass sheet degeneracy \citep[e.g.,][]{falco1985,gorenstein1988,saha2000,schneider2013,xu2016}, a measurement of the lens galaxy's velocity dispersion \citep[e.g.,][]{koopmans2003,koopmans2004} and an estimate of the external convergence, $\kext$, along the line of sight (LOS) is needed. $\kext$ can also bias the lens model parameters if unaccounted for \citep[e.g.,][]{collett2013,greene2013,mccully2014,mccully2016}. In an effort to provide an accurate independent estimate of $H_{0}$ using time-delay lenses, we use a number of new datasets as part of our project, $H_{0}$ Lenses in COSMOGRAIL's Wellspring (H0LiCOW), to model five lensed quasars. These datasets include high-resolution imaging with the {\it Hubble Space Telescope} (\hst), precise time-delay measurements from the COSmological MOnitoring of GRAvItational Lenses \citep[COSMOGRAIL;][]{courbin2005,eigenbrod2005,bonvin2016b} project and from Very Large Array (VLA) monitoring \citep{fassnacht2002}, a photometric and spectroscopic survey to characterize the LOS mass distribution to estimate $\kext$ in these systems, and stellar velocity dispersion measurements of the strong lens galaxies. With five separate lenses, we plan to account for systematic uncertainties and obtain a robust constraint on $H_{0}$ to $< \combprec\%$ precision. In this paper, we present the results of a detailed lens modeling analysis of the gravitational lens \hequad~using new high-resolution imaging data from \hst. \hequad~is the third H0LiCOW system analyzed in this manner, following \blens~\citep{suyu2010b} and \rxjlens~\citep{suyu2013, suyu2014}. This paper is the fourth in a series of papers detailing our analysis of \hequad. The other papers include an overview of the H0LiCOW project \citep[][hereafter H0LiCOW Paper I]{suyu2016}, a spectroscopic survey of the \hequad~field and a characterization of the groups along the LOS \citep[][hereafter H0LiCOW Paper II]{sluse2016}, a photometric survey of the \hequad~field and an estimate of $\kext$ due to the external LOS structure \citep[][hereafter H0LiCOW Paper III]{rusu2016}, and a presentation of our latest time-delay measurements for \hequad~and the cosmological inference from our combined analysis of \hequad, \blens, and \rxjlens~\citep[][hereafter H0LiCOW Paper V]{bonvin2016a}. This paper is organized as follows. We provide a brief overview of using time-delay lenses for cosmography in \sref{sec:theory}. In \sref{sec:data}, we describe the observational data used in our analysis. We describe our lens modeling procedure in \sref{sec:lensmod}. The time-delay distance results and their implications for cosmology are presented in \sref{sec:results}. We summarize our main conclusions in \sref{sec:conclusions}. Throughout this paper, all magnitudes given are on the AB system.

1607.01403

1607

1607.04226_arXiv.txt

Modeling of gravitational waves (GWs) from binary black hole inspiral brings together early post-Newtonian waveforms and late quasinormal ringing waveforms. Attempts to bridge the two limits without recourse to numerical relativity involve predicting the time of the peak GW amplitude. This prediction will require solving the question of why the peak of the ``source,'' i.e., the peak of the binary angular velocity, does not correspond to the peak of the GW amplitude. We show here that this offset can be understood as due to the existence of two distinct components of the radiation: the ``direct'' radiation analogous to that in flat spacetime, and ``scattered'' radiation associated with curved spacetime. The time dependence of these two components, and of their relative phases determines the location of the peak amplitude. We use a highly simplified model to clarify the two-component nature of the source, then demonstrate that the explanation is valid also for an extreme mass ratio binary inspiral.

\label{sec:overview} \subsection{Introduction and Summary} \medskip The recent detections of gravitational radiation events GW150914~\cite{GWPRL1} and GW151226~\cite{GWPRL2} from black hole binary inspiral has underscored the importance of understanding the theoretical predictions of the waveforms for such processes. Numerical relativity, the computation of the Einstein's nonlinear equations~\cite{NR}, has played an important role in generating these predictions, but there is a wide range of binary parameters, and studying an appropriately large and finely spaced set of waveforms requires a more efficient methodology to supplement numerical relativity. One such methodology is the effective-one-body approximation (EOB)~\cite{EOB1,EOB2,BB}, used together with the linear equations of the extreme mass ratio particle-perturbation technique~\cite{BBHKOP,TBKH} and numerical relativity~\cite{PBBBKPS}. In the EOB model, two different techniques are used to generate a full waveform signal from a black hole binary system and the results are amplitude and phase matched. In the early epoch (inspiral-plunge) of the binary's evolution, an EOB Hamiltonian system is evolved and the waveform is computed using an inspiral-plunge trajectory. For the late (merger-ringdown) epoch, the early epoch waveform is matched to a linear superposition of several quasi-normal modes. The precise moment in the binary's evolution at which this matching is performed is thus a very critical matter in the EOB approach. In the original EOB work~\cite{EOB1,EOB2} it was argued that this matching should be performed close to the light ring where the orbital frequency peaks; thus, the peak of the waveform amplitude was identified with the light ring as well. The fact that these two peaks do not actually align has been a complication in this program. Particularly troublesome has been the dependence of the peak offsets on the orbital frequency, and the fact that different (spin-weighted) spherical harmonic modes have significantly different offsets~\cite{BBHKOP,TBKH,PBBBKPS}. In the EOB papers a time delay is defined, $\Delta t^{\ell,m}_{\rm peak}=t^{\ell,m}_{\rm peak} -t^{\Omega}_{\rm peak} $, where $t^{\ell,m}_{\rm peak}$ is the time at which the peak of the ${\ell, m}$ mode of radiation is observed, and $t^{\Omega}_{\rm peak}$ is the time at which radiation would be observed from the maximum of the particle's angular velocity. In most cases (i.e. results from numerical relativity and particle perturbation theory models), this delay is negative; the peak radiation is earlier than the radiation from the peak of angular velocity. Such studies have raised the question why there should be a peak offset, that is, why the apparent peak of the supposed source of radiation does not align with the peak of the radiation generated. Not only would an answer satisfy scientific curiosity, but it could lead to an effective way of predicting the peak offsets. It has been suggested~\cite{ScottField} that the offset, and its dependence on angular frequency, might be explained as manifestations of the energy dependence of the scattering of gravitational radiation off the curvature potential~\cite{curvpot} created by the black holes. In this paper we will answer that question by exploiting an approach we have recently developed~\cite{PNK}. In that approach, we use the Fourier domain Green function (FDGF) for a very simple model in which we replace the curvature potential with a simplified truncated dipole potential (TDP)~\cite{tdp}. In our approach, we use point particle trajectories, but we do not require them to be analogs of geodesics. This greater freedom to vary trajectories, combined with the simplification from the TDP model, gives transparency to the problem of the location of the radiation peak, and its dependence on angular velocity. The clear picture that is revealed is that the outgoing radiation consists of two components. One component, the quasinormal or scattered radiation, is due to the motion through the strong field spacetime (equivalently, motion in the neighborhood of the peak of the curvature potential). The second component is the ``direct radiation,'' radiation that has little to do with the curved spacetime, and can be fully ascribed to the motion of the particle as if it were in flat spacetime. The peak radiation is a result of the combination of the two components. Through the use of numerical evolution codes we have investigated whether this picture applies to the extreme mass-ratio inspiral (EMRI) problem in the Schwarzschild geometry. Although the direct and scattered contributions cannot be separated in an evolution code, the character of the results clearly indicates that this picture does apply. Further study is appropriate to the applicability to the inspiral of comparable mass holes, but the indications are strong that here too it applies. The remainder of this paper is organized as follows. In Sec.~\ref{sec:TDP} we review the features of the TDP/FDGF model that are most relevant to the investigation of the peak location. Numerical results from that model are presented in Sec.~\ref{sec:TDPresults}, and a tentative explanation is presented of the location of the amplitude peak. Section~\ref{sec:Schw} then shows that the qualitative features of the peak location are the same for the EMRI models in the Schwarzschild spacetime as in the TDP model. Conclusions are given in Sec.~\ref{sec:conc}.

\label{sec:conc} Though peak time offsets have been important in waveform construction, all that has been known is that such a time offset exists and that it is different (a) in the Schwarzschild background for gravitational waves in different multipole modes, and also (b) in the Kerr background for even a given multipole mode, but different values of $a/M$. We know of no claim in the literature that suggests an explanation of this offset, what it depends on, or whether (a) and (b) are related. In this paper we present insights that provide at least the beginning of an explanation of the peak offset phenomenon. These insights are the results of exploiting the simplicity of a model first explored in our previous paper~\cite{PNK}. In particular, (i)~the curvature potential of the Schwarzschild EMRI problem is replaced by a simplified truncated dipole potential; (ii)~the source is understood through the Fourier domain Green function; (iii)~we do not restrict ourselves to geodesic orbits, but use trajectories that allow us to probe the mechanisms for the generation of radiation. This simplicity has produced a very simple explanation of why the amplitude peak for low $\omax$ is at the ``expected'' late retarded time, but for high $\omax$ is significantly earlier. In simplest terms, the idea is that the radiation consists of two components, direct radiation and scattered radiation, and that the direct component vanishes at the ``expected'' retarded time. For low $\omax$ the two components are out of phase, so the late-time vanishing of the subtracted direct radiation means that the peak, dominated by the scattered radiation, can occur at late times. For high $\omax$ the two components are in phase. The vanishing of the late-time direct radiation, therefore, removes its addition to the total, and pushes the peak total to earlier times. The separation of the direct and scattered components of the radiation for the Schwarzschild EMRI problem is not simple, but the qualitative similarity of our results for the TDP and for the Schwarzschild background, both for gravitational waves and for scalar waves (not reported here), makes a strong case that the same mechanism is at work. The case is particularly strong for the claim that the replacement of the Schwarzschild curvature potential by the TDP potential is justified in seeking an explanation. Further support comes from the details of the Schwarzschild results in the literature. In Table II of the EMRI studies reported for the Schwarzschild background in Ref.~\cite{BBHKOP}, the time by which the amplitude peak precedes the angular velocity peak, for a given $\ell$ mode, increases with the $m$ index of the mode. This increase is mathematically equivalent to increasing the angular velocity. The reported results, then, show that as the angular velocity increases, the peak moves earlier, as in the results we report above, both for the TDP and Schwarzschild models. Turning to the Kerr case, in Table \ref{Kerrtable} we present results \begin{table} \begin{tabular}{|c|c|c|c|} \hline $a/M$& $h_{22}$\ offset& $2M\omega_{\rm peak}$& $M\omega_{\rm QN}$ \\ \hline -0.9 & +2.0 & 0.18 & 0.30 \\ 0.0 & -3.0 & 0.28 & 0.37 \\ 0.5 & -7.0 & 0.38 & 0.47 \\ 0.9 & -40.0 & 0.64 & 0.67 \\ 0.99 & -60 & 0.84 & 0.87 \\ \hline \end{tabular} \caption{The time offsets of the $\ell,m=2,2$ peak of radiation from particles spiralling into Kerr black holes with spin parameter $a$, and mass $M$. The second column shows the offset (in units of $M$) by which the GW peak follows the signal from the maximum of the angular velocity. The third column shows twice the peak orbital frequency (in units of $1/M$) of the inspiralling particle. The fourth column is the quasinormal frequency of the dominant (least damped) $\ell,m=2,2$ mode. It is clear that when the last two columns have values that are comparable, a significant offset is observed. \label{Kerrtable}} \end{table} of EMRI studies adapted from Refs.~\cite{BBHKOP,TBKH}. Of note is the fact that the peak frequency is at a value quite close to that of the real part of the least damped quasinormal frequency for the GW 2,2 mode. These frequencies correspond to the high/intermediate frequencies of Secs.~\ref{sec:TDPresults} and \ref{sec:Schw} and this strongly suggests that the same mechanism may be at play in the Kerr case as well. It remains to be seen whether the strong $a/M$ dependence of the offsets in the Kerr geometry can be fully explained in terms of the same mechanism as in the TDP and Schwarzschild cases. The insights about these peak offsets, as presented in this paper, can most effectively be used in the EOB approach to waveform construction.

1607.04226

1607

1607.06479_arXiv.txt

We examine the dark matter content of satellite galaxies in $\Lambda$CDM cosmological hydrodynamical simulations of the Local Group from the APOSTLE project. We find excellent agreement between simulation results and estimates for the $9$ brightest Galactic dwarf spheroidals (dSphs) derived from their stellar velocity dispersions and half-light radii. Tidal stripping plays an important role by gradually removing dark matter from the outside in, affecting in particular fainter satellites and systems of larger-than-average size for their luminosity. Our models suggest that tides have significantly reduced the dark matter content of Can~Ven~I, Sextans, Carina, and Fornax, a prediction that may be tested by comparing them with field galaxies of matching luminosity {\it and} size. Uncertainties in observational estimates of the dark matter content of individual dwarfs have been underestimated in the past, at times substantially. We use our improved estimates to revisit the `too-big-to-fail' problem highlighted in earlier N-body work. We reinforce and extend our previous conclusion that the APOSTLE simulations show no sign of this problem. The resolution does {\it not} require `cores' in the dark mass profiles, but, rather, relies on revising assumptions and uncertainties in the interpretation of observational data and accounting for `baryon effects' in the theoretical modelling.

\label{SecIntro} The steep slope of the dark matter halo mass function at the low-mass end is a defining characteristic of the $\Lambda$CDM cosmological paradigm. It is much steeper than the faint-end slope of the galaxy stellar mass function, implying that low-mass CDM haloes are significantly more abundant than faint dwarf galaxies \citep{Moore1999b,Klypin1999}. This discrepancy is usually reconciled by assuming that dwarfs form preferentially in relatively massive haloes, because cosmic reionization and the energetic feedback from stellar evolution are effective at removing baryons from the shallow gravitational potential of low-mass systems and at curtailing their star forming activity \citep{Bullock2000,Benson2002,Somerville2002}. Such scenario makes clear predictions for the stellar mass -- halo mass relation at the faint end. A simple -- but powerful and widely used -- parameterization of this prediction is obtained from abundance matching (AM) modeling, where galaxies and CDM haloes are ranked by mass and matched to each other respecting their relative ranked order \citep[see, e.g.,][]{Frenk1988,Vale2004,Guo2011,Moster2013,Behroozi2013}. Halo masses may thus be derived from stellar masses, yielding clear predictions amenable to observational testing. Most such tests rely on using kinematic tracers such as rotation speeds or velocity dispersions to estimate the total gravitational mass enclosed within the luminous radius of a galaxy. Its dark matter content, computed after subtracting the contribution of the baryons, may then be used to estimate the total virial\footnote {We define virial quantities as those corresponding to the radius where the spherical mean density equals 200 times the critical density for closure, $3H^2/8\pi G$. Virial quantities are identified by a ``200'' subscript.} mass of the system. Such estimates rely heavily on the similarity of the mass profiles of CDM haloes \citep[][referred to hereafter as NFW]{Navarro1996,Navarro1997}, and involve a fairly large extrapolation, since virial radii are typically much larger than galaxy radii. These tests have revealed some tension between the predictions of AM models and observations. \citet{Boylan-Kolchin2011}, for example, estimated masses for the most luminous Galactic satellites that were lower than those of the most massive substructure haloes in N-body simulations of Milky Way-sized haloes from the Aquarius Project \citep{Springel2008b}. \citet{Ferrero2012} reported a related finding when analyzing the dark matter content of faint dwarf irregular galaxies in the field: many of them implied total virial masses well below those predicted by AM models. Subsequent work has highlighted similar results both in the analysis of M31 satellites \citep{Tollerud2014,Collins2014}, as well as in other samples of field galaxies \citep{Garrison-Kimmel2014,Papastergis2015}. These discrepancies may in principle be reconciled with $\Lambda$CDM in a number of ways. One possibility is to reconsider virial mass estimates based on the dark mass enclosed by the galaxy, a procedure that is highly sensitive to assumptions about the halo mass profile. A popular revision assumes that the assembly of the galaxy may lead to a reshuffling of the mass profile, pushing dark matter out of the inner regions and creating a constant-density `core' in an otherwise cuspy NFW halo \citep[e.g.,][]{Navarro1996b,Mashchenko2006,Governato2012}. These cores allow dwarf galaxies to inhabit massive haloes despite their relatively low inner dark matter content. This option has received some support from hydrodynamical simulations \citep[see, e.g.,][for a review]{Pontzen2014} although the results are sensitive to how star formation and feedback are implemented. Indeed, no consensus has yet been reached over the magnitude of the effect, its dependence on mass, or even whether such cores exist at all \citep[see, e.g.,][and references therein]{Parry2012,Garrison-Kimmel2013,DiCintio2014,Schaller2015a,Oman2015, Onorbe2015}. A second possibility is that Galactic satellites have been affected by tidal stripping, which would preferentially remove dark matter \citep[e.g.,][]{Penarrubia2008b} and, therefore, act to reduce their dark mass content, much as the baryon-induced `cores' discussed in the preceding paragraph. This proposal would not help to solve the issue raised by field dwarf irregulars \citep{Ferrero2012} nor the low dark matter content of Galactic satellites (tides are, of course, already included in N-body halo simulations), unless baryon-induced cores help to enhance the effects of tides, as proposed by \citet{Zolotov2012} and \citet{Brooks2014}. A third option is to revise the abundance matching prescription so as to allow dwarf galaxies to inhabit haloes of lower mass. This would be the case if some galaxies simply fail to form (or are too faint to feature in current surveys) in haloes below some mass: once these ``dark'' systems are taken into account, the AM stellar mass -- halo mass relation would shift to systematically lower virial masses for given stellar mass, as pointed out by \citet{Sawala2013}. The existence of `dark' subhaloes does not, on its own, solve the problem pointed out by \citet{Boylan-Kolchin2011}, which is usually referred to as the `too-big-to-fail' problem \citep[hereafter TBTF, see also][]{Boylan-Kolchin2012}. Indeed, associating dwarfs with lower halo masses would not explain why many of the most massive substructures in the Aquarius haloes seem inconsistent with the kinematic constraints of the known Galactic satellites. One explanation might be that fewer massive subhaloes are present in the Milky Way (MW) than in Aquarius haloes. Since the number of substructures scales with the virial mass of the main halo, a lower Milky Way mass would reduce the number of massive substructures, thus alleviating the problem \citep{Wang2012a,Vera-Ciro2013,Cautun2014}. Another possibility is that dark-matter-only (DMO) simulations like Aquarius overestimate the subhalo mass function. Low mass haloes are expected to lose most of their baryons to cosmic reionization and feedback, a loss that stunts their growth and reduces their final mass. The effect is limited in terms of mass (baryons, after all, make up only $17$ per cent of the total mass of a halo) but it can have disproportionate consequences on the number of massive substructures given the steepness of the subhalo mass function \citep{Guo2015,Sawala2016}. \begin{figure} \hspace{-0.2cm} \resizebox{8cm}{!}{\includegraphics{figures/Fig1.ps}}\\% \caption{{\it Top}: Stellar velocity dispersion and effective radius ($R_{\rm eff}$) of the Fornax dSph. The $R_{\rm eff}$ distribution is obtained by convolving uncertainties in distance and in the observed angular half-light radius, using uncertainties from the literature and assuming Gaussian error distributions. {\it Middle}: Dynamical mass within the deprojected 3D half-light radius ($r_{1/2}$) of Fornax, calculated using eq.~\ref{eqW10} \citep{Wolf2010}. The red histogram shows the result of propagating the observational uncertainties, whereas the grey histogram adds a $23$ per cent base modeling uncertainty, as suggested by \citet{Campbell2016}. {\it Bottom}: Circular velocity at $r_{1/2}$ ($V_{1/2}$), including both observational and systematic uncertainties, calculated from the final $M_{1/2}$ distribution (middle panel). Unlike $M_{1/2}$, $V_{1/2}$ is independent of $r_{1/2}$. Contours in all panels enclose $50$ per cent and $80$ per cent of the distributions. } \label{FigError} \end{figure} \begin{figure} \hspace{-0.2cm} \resizebox{8cm}{!}{\includegraphics{figures/Fig2.ps}}\\% \caption{{\it Top:} Circular velocity at the half-light radius of Milky Way classical dSphs. Open circles show the results of \citet{Wolf2010} with $1\sigma$ error. The bar-and-whisker symbols show the results of this work, including observational and systematic uncertainties (see, e.g., the bottom panel of Fig.~\ref{FigError} for the case of the Fornax dSph). The thick and thin portions illustrate interquartile and $10$--$90^{\rm th}$ percentile intervals, respectively. Our results suggest that $V_{1/2}$ uncertainties have been underestimated in previous work. Slanted lines show objects with constant crossing time, as labelled. {\it Bottom:} Stellar mass derived for the $9$ Galactic dSphs, shown as a function of their half-light radius. The blue dashed line indicates the characteristic halo mass -- radius dependence of APOSTLE centrals, computed from the fit shown in Fig.~\ref{FigMstarM200}. The line divides the sample in two groups of compact objects resilient to tides and more extended systems where tidal effects may be more apparent.} \label{FigVR} \end{figure} We explore these issues here using $\Lambda$CDM cosmological hydrodynamical simulations of the Local Group from the APOSTLE\footnote{A Project Of Simulating The Local Environment} project \citep{Fattahi2016,Sawala2016}. These simulations use the same code as the EAGLE project, whose numerical parameters have been calibrated to reproduce the galaxy stellar mass function and the distribution of galaxy sizes \citep{Schaye2015,Crain2015}. Our analysis complements that of \citet{Sawala2016}, who showed that APOSTLE reproduces the Galactic satellite luminosity/stellar mass function, as well as the total number of galaxies brighter than $10^5 \Msun$ within the Local Group. We extend here the TBTF discussion of that paper by reviewing the accuracy of observational constraints (Sec.~\ref{SecObs}), which are based primarily on measurements of line-of-sight velocity dispersions and the stellar half-mass radii ($r_{1/2}$) of `classical' (i.e., $M_V<-8$) Galactic dwarf spheroidals (dSphs), and by focusing our analysis on the actual mass enclosed within $r_{1/2}$ rather than on extrapolated quantitites such as the maximum circular velocity of their surrounding haloes. We also highlight the effect of Galactic tides, and identify the satellites where such effects might be more easily detectable observationally. This paper is organized as follows. We begin by reviewing in Sec.~\ref{SecObs} the observational constraints on the mass of Galactic dSphs. We then describe briefly our simulations and discuss our main results in Sec.~\ref{SecResults}, and conclude with a summary of our main conclusions in Sec.~\ref{SecConc}.

\label{SecConc} We use the APOSTLE suite of $\Lambda$CDM cosmological hydrodynamical simulations of the Local Group to examine the masses of satellite galaxies brighter than $M_V=-8$ (i.e., $M_{\rm str}>10^5 \Msun$). Our analysis extends that of \citet{Sawala2016}, were we showed that our simulations reproduce the Galactic satellite luminosity function and show no sign of either the `missing satellites' problem nor of the `too-big-to-fail' problem highlighted in earlier work. Our main conclusions may be summarised as follows. Previous studies have underestimated the uncertainty in the mass enclosed within the half-light radii of Galactic dSphs, derived from their line-of-sight velocity dispersion and half-light radii. Our analysis takes into account the error propagation due to uncertainties in the distance, effective radius, and velocity dispersion, and also include an estimate of the intrinsic dispersion of the modeling procedure, following the recent work of \citet{Campbell2016}. The latter is important as it introduces a base systematic uncertainty that exceeds $\sim 20$ per cent. Simulated galaxies in APOSTLE/EAGLE follow a stellar mass -- halo mass relation that differs, for dwarf galaxies, from common extrapolations of abundance matching models, a difference that is even more pronounced for satellites due to tidal stripping, At fixed stellar mass, APOSTLE dwarfs inhabit halos significantly less massive than AM predicts. This difference, however, might not be readily apparent because tides strip halos from the outside in and some dSphs are too compact for tidal effects to be readily apparent. We find that the dynamical mass of {\it all} Galactic dSphs is in excellent agreement with that of APOSTLE satellites that match their stellar mass. APOSTLE centrals (i.e., not satellites), on the other hand, overestimate the observed mass of four Galactic dSphs (Can Ven~I, Sextans, Carina, and Fornax), suggesting that they have had their dark matter content significantly reduced by stripping. The other, more compact, dSphs are well fit by either APOSTLE satellites or centrals, so tides are not needed to explain their dark matter content. After accounting for tidal mass losses, we find that all APOSTLE halos (satellites or centrals) with $V_{\rm max}>25 \kms$ host dwarfs brighter than $M_V=-8$. Only systems with fewer than $\sim 12$ subhaloes with $V_{\rm max}>25$ km/s are thus compatible with the population of luminous MW satellites. This suggests an upper limit to the mass of the Milky Way halo: we find that most halos with virial mass not exceeding $2\times 10^{12}\, M_\odot$ should pass this test, unless they are unusually overabundant in massive substructures. Our APOSTLE primaries satisfy these constraints, and show a dwarf galaxy population in agreement with observations of the Local Group, including their abundance as a function of mass, their dark matter content, and their global kinematics. Furthermore, APOSTLE uses the same galaxy formation model that was found by EAGLE to reproduce the galaxy stellar mass function in cosmologically significant volumes. We consider this a significant success for direct simulations of galaxy formation based on the $\Lambda$CDM paradigm. We note that this success does not require any substantial modification to well-established properties of $\Lambda$CDM. In particular, none of our simulated dwarf galaxies have `cores' in their dark mass profiles, but yet have no trouble reproducing the detailed properties of Galactic satellites. Baryon-induced cores are not mandatory to solve the `too-big-to-fail' problem. We end by noting that a number of recent studies have argued that TBTF-like problems also arise when considering the properties of M31 satellites \citep{Tollerud2014,Collins2014}, as well as those of field galaxies in the local Universe \citep{Garrison-Kimmel2014,Papastergis2015}. It remains to be seen whether the resolution we advocate here for Galactic satellites will solve those problems as well. We plan to report on those issues in future work.

1607.06479

1607

1607.03540_arXiv.txt

The merger of binary neutron-stars systems combines in a single process: extreme gravity, copious emission of gravitational waves, complex microphysics, and electromagnetic processes that can lead to astrophysical signatures observable at the largest redshifts. We review here the recent progress in understanding what could be considered Einstein's richest laboratory, highlighting in particular the numerous significant advances of the last decade. Although special attention is paid to the status of models, techniques, and results for fully general-relativistic dynamical simulations, a review is also offered on initial data and advanced simulations with approximate treatments of gravity. Finally, we review the considerable amount of work carried out on the post-merger phase, including: black-hole formation, torus accretion onto the merged compact object, connection with gamma-ray burst engines, ejected material, and its nucleosynthesis.

\label{sec:Introduction} Neutron stars are believed to be born in supernova explosions triggered by the collapse of the iron core in massive stars. Many astronomical observations have revealed that binary neutron stars\footnote{\lrn{With \textit{``binary neutron-star''} systems we here refer to binary systems composed of two neutron stars; in astronomy, such systems are often called \textit{``double neutron-star''} systems in order to distinguish them from binary systems in which one star is a neutron star and the other a white dwarf.}} (BNSs) indeed exist \cite{Kramer04} and the most important physical properties of all known such systems are collected in Table \ref{tab:observedNS}. Despite this observational evidence of existence, the formation mechanisms of BNS systems are not known in detail. The general picture is that in a binary system made of two massive main-sequence stars \lrn{of masses between approximately $8$ and $25\,M_\odot$}, the more massive one undergoes a supernova explosion and becomes a neutron star. This is followed by a very uncertain phase in which the neutron star and the main-sequence star evolve in a ``common envelope'', that is, with the neutron star orbiting in the extended outer layers of the secondary star \cite{Kiziltan:2010a,Ivanova2013,Ozel2016}. At the end of this stage, also the second main-sequence star undergoes a supernova explosion and, if the stars are still bound after the explosions, a BNS system is formed. The common-envelope phase, though brief, is crucial because in that phase the distance between the stars becomes much smaller as a result of drag, and this allows the birth of BNS systems that are compact enough to merge within a Hubble time, following the dissipation of their angular momentum through the emission of gravitational radiation. \lrn{It is also possible that during the common-envelope phase the neutron star collapses to a black hole, thus preventing the formation of a BNS.} Another possible channel for the formation of BNS systems may be the interaction of two isolated neutron stars in dense stellar regions, such as globular clusters, in a process called ``dynamical capture'' \cite{Oleary2009, Lee2010, Thompson2011}. Dynamically formed binary systems are different from the others because they have higher ellipticities (see Sect. \ref{sec:dynamical-capture}). It is presently not known what fraction of BNS systems would originate from dynamical capture, but it is expected that these binaries are only a small part of the whole population. \begin{table*}[ht] \caption{\label{tab:observedNS} Observational data of neutron stars in binary neutron-star systems containing a pulsar. Reported in the various columns are: the name of the binary, the total (gravitational) mass $M_{\rm tot}$, the (gravitational) masses of the pulsar and that of its neutron-star companion $M_{{\rm A}}, M_{_{\rm B}}$, the mass ratio $q \leq 1$, the orbital period $T_{\rm orb}$, the projected semi-major axis of the orbit $R$ (\ie the projection of the semi-major axis onto the line of sight), the orbital eccentricity $e_{\rm orb}$, the distance from the Earth, the barycentric rotation frequency $f_{\rm s}$, and the inferred surface magnetic dipole field $B_{\rm surf}$. The data are taken from the respective references and truncated to four significant digits for the masses and to two significant digits for the rest. Note that in the case of the "double pulsar" system J0737-3039 (the only double system where both neutron stars are detectable as pulsars), the magnetic field of the second-formed pulsar (not reported in this table) is estimated to be 1.59E+12 G.} \begin{tabular}[t]{lrrrrrrrrrr} \hline \hline Name & $M_{\rm tot}$ & $M_{_{\rm A}}$ & $M_{_{\rm B}}$ & $q$ & $T_{\rm orb}$ & $R$ & $e_{\rm orb}$ & $D$ & $f_{\rm s}$ & $B_{\rm surf}$ \\ [0.5ex] & $[{M_{\odot}}]$ & $[{M_{\odot}}]$ & $[{M_{\odot}}]$ & & [days] & [light s] & & [kpc] & [Hz] & [G] \\ [0.5ex] \hline J0453+1559~\cite{Martinez2015} & 2.734 & 1.559 & 1.174 & 0.75 & 4.1 & 14 & 0.11 & 1.8 & 22 & 9.3E+09 \\ [0.5ex] J0737-3039~\cite{Kramer2006}~~~~~~~~~~ & 2.587 & 1.338 & 1.249 & 0.93 & 0.10 & 1.4 & 0.088 & 1.1 & 44 & ~~~6.4E+09 \\ [0.5ex] J1518+4904~\cite{Janssen2008} & 2.718 & $\,\,\,\, <$1.766 & $\,\,\,\, >$0.951 & $\,\,\,\, >$0.54 & 8.6 & 20 & 0.25 & 0.7 & 24 & 9.6E+08 \\ [0.5ex] B1534+12~\cite{Fonseca2014} & 2.678 & 1.333 & 1.345 & 0.99 & 0.42 & 3.7 & 0.27 & 1.0 & 26 & 9.6E+09 \\ [0.5ex] J1753-2240~\cite{Keith2009} & -- & -- & -- & -- & 14 & 18 & 0.30 & 3.5 & 10 & 9.7E+09 \\ [0.5ex] J1756-2251~\cite{Ferdman2014} & 2.577 & 1.341 & 1.23 & 0.92 & 0.32 & 2.8 & 0.18 & 0.73 & 35 & 5.4E+09 \\ [0.5ex] J1807-2500B~\cite{Lynch2012} & 2.571 & 1.366 & 1.21 & 0.89 & 1.0 & 29 & 0.75 & -- & 239 & $\,\,\,\, \leq$9.8E+08 \\ [0.5ex] J1811-1736~\cite{Corongiu2007} & 2.571 & $\,\,\,\, <$1.478 & $\,\,\,\, >$1.002 & $\,\,\,\, >$0.68 & 19 & 35 & 0.83 & 5.9 & 9.6 & 9.8E+09 \\ [0.5ex] J1829+2456~\cite{Champion2004} & 2.59\phantom{0} & $\,\,\,\, <$1.298 & $\,\,\,\, >$1.273 & $\,\,\,\, >$0.98 & 1.2 & 7.2 & 0.14 & 0.74 & 24 & 1.5E+09 \\ [0.5ex] J1906+0746~\cite{vanLeeuwen2015} & 2.613 & 1.291 & 1.322 & 0.98 & 0.17 & 1.4 & 0.085 & 7.4 & 6.9 & 1.7E+12 \\ [0.5ex] J1913+1102~\cite{Lazarus2016} & 2.875 & $\,\,\,\, <$1.84\phantom{0} & $\,\,\,\, >$1.04\phantom{0} & $\,\,\,\, >$0.56 & 0.21 & 1.8 & 0.090 & 13 & 1.1 & 2.1E+09 \\ [0.5ex] B1913+16~\cite{Weisberg2016} & 2.828 & 1.449 & 1.389 & 0.96 & 0.32 & 2.3 & 0.62 & 7.1 & 17 & 2.3E+10 \\ [0.5ex] J1930-1852~\cite{Swiggum2015} & 2.59\phantom{0} & $\,\,\,\, <$1.199 & $\,\,\,\, >$1.363 & $\,\,\,\, >$0.88 & 45 & 87 & 0.40 & 2.3 & 5.4 & 6.0E+10 \\ [0.5ex] B2127+11C~\cite{Jacoby2006} & 2.713 & 1.358 & 1.354 & 1.0\phantom{0} & 0.34 & 2.5 & 0.68 & 13 & 33 & 1.2E+10 \\ [0.5ex] \hline \hline \end{tabular} \end{table*} This is undoubtedly an exciting and dynamical time for research on BNS mergers, when many accomplishments have been achieved (especially since 2008), while many more need to be achieved in order to describe such fascinating objects and the related physical phenomena. The first direct detection through the advanced interferometric LIGO detectors \cite{Harry2010} of the gravitational-wave signal from what has been interpreted as the inspiral, merger and ringdown of a binary system of black holes \cite{Abbott2016a} marks, in many respects, the beginning of gravitational-wave \lrn{astronomy; a second detection was made a few months later \cite{Abbot2016g}}. Additional advanced detectors, such as Virgo \cite{Accadia2011_etal}, KAGRA \cite{Aso:2013} \lrn{and LIGO India (see \eg \cite{Fairhurst2014})}, are going to become operational in the next few years, and we are likely to witness soon also signals from the inspiral and post-merger of neutron-star binaries or neutron-star--black-hole binaries, with a detection rate that has an uncertainty of three orders of magnitude, but is expected to be of several events per year \cite{Abadie:2010_etal}. BNS mergers are rather unique objects in the landscape of relativistic astrophysics as they are expected to be at the origin of several and diverse physical processes, namely: {(i)} to be significant sources of gravitational radiation, not only during the inspiral, but also during and after the merger; {(ii)} to be possible progenitors for short-gamma-ray bursts (SGRBs); {(iii)} to be the possible sources of other electromagnetic and neutrino emission; {(iv)} to be responsible for the production of a good portion of the very heavy elements in the Universe. When viewed in this light, BNS mergers naturally appear as Einstein's richest laboratory, where highly nonlinear gravitational effects blend with complex microphysical processes and yield astonishing astrophysical phenomena. As we will discuss in more detail in the following Section, the typical scenario leading to SGRBs assumes that a system composed of a rotating black hole and a surrounding massive torus is formed after the merger \cite{Narayan92,Eichler89}. A large number of numerical simulations \cite{Shibata99d, Baiotti08, Anderson2007, Liu:2008xy, Bernuzzi2011} have confirmed that this scenario can be attained through BNS mergers unless the progenitor stars have very small masses [smaller than half of the maximum allowed mass for neutron stars with a given equation of state (EOS)], \lrn{or when the merged object collapses to a black hole as a uniformly rotating neutron star in vacuum \cite{Margalit2015}}. Furthermore, if sufficiently massive, the torus could provide the large amount of energy observed in SGRBs, either through neutrino processes or by extracting the rotational energy of the black hole via magnetic fields \cite{Paczynski86, Eichler89}. Furthermore, if the neutron stars in the binary have relatively large magnetic fields and extended magnetospheres, the inspiral could also be accompanied by a precursor electromagnetic signal \cite{Palenzuela2013a}, while after the merger magnetically confined jet structures may form once a torus is present around the black hole \cite{Rezzolla:2011, Paschalidis2014, Dionysopoulou2015, Ruiz2016}. Possible evidence that a BNS merger can be behind the phenomenology associated with SGRBs has emerged recently from the infrared excess in the afterglow curve of Swift's short gamma-ray burst SGRB 130603B \cite{Berger2013, Tanvir2013}, which has been interpreted as a ``macronova'' emission \cite{Li1998, Kulkarni2005_macronova-term} (sometimes also referred to as ``kilonova'' \cite{Metzger:2010}), \ie as due to the radioactive decay of by-products of the $r$-processed matter from the material ejected in the merger\footnote{We will discuss this further in Section \ref{sec:bph}, but we briefly recall here that $r$ (or rapid) processes are nucleosynthetic processes involving the rapid capture of neutrons.}. Other macronova candidates, \lrn{\eg GRB 060614 and GRB 050709 \cite{Yang2015,Jin2016},} are presently being considered. For instance, strong evidence for a macronova component has been found recently in the peculiar long-short event GRB 060614 \cite{Yang2015} and in its afterglow \cite{Jin2015}, while a careful re-examination of the afterglow of SGRB 050709, the first short event with an identified optical afterglow, has highlighted a macronova component \cite{Jin2016}. The observations of the infrared transient in these afterglows are important not only because they provide a potential observational link between two distinct phenomena (\ie a SGRB explosion and a radioactive decay), but also because they suggest that BNSs can be the site of active and intense nucleosynthesis. Additional evidence in this direction is offered by the Solar system abundance of $^{244}$Pu \cite{wallner:15, Hotokezaka:2015b} and recent observations of $r$-process enriched stars in a metal-poor ultra-faint dwarf galaxy \cite{Ji:15}. Both of these observations suggest that $r$-process elements might be preferentially produced in rare/high-yield events such as mergers instead of common/low-yield occurrences such as core-collapse supernovae. Given the complex nonlinear nature of merging BNSs, it is inevitable that fully three-dimensional numerical simulations are the only tool available for studying these processes accurately and with a sufficient degree of realism. At present, there are about a dozen numerical codes in groups across the world that are able to produce meaningful results about BNS mergers. Most of these codes solve the full Einstein equations without approximations, together with the equations of relativistic hydrodynamics and/or (resistive) magnetohydrodynamics (MHD) equations. However, there are also codes that treat matter with smoothed-particle-hydrodynamics (SPH) methods and with some approximate treatment of gravity, which is however balanced by more advanced treatments in the microphysical processes. Each of these codes represents a complex computational infrastructure built over the last decade (if not more) and that in most cases already provides, together with an accurate description of the bulk motion of matter (before and after the merger), also an approximate representation of the microphysical aspects related to the EOS, to the neutrino radiation transport, to the nuclear reactions taking place in the ejected matter, and, ultimately, to the electromagnetic signal from merging BNSs. Such computational infrastructures are being continuously updated and improved, either through the use of more advanced numerical methods, through the development of novel formulations of the equations, or through the introduction of new and more refined levels of microphysical description. Finally, all of these codes also share common scientific goals: a faithful representation of the gravitational-wave signal produced before and after the merger, as well as an interpretative and predictive description of the phenomenology behind SGRBs. This Report is meant to provide a general but possibly detailed description of the progress achieved in the numerous areas touched up by investigations of BNS mergers and hence to provide a snapshot of the status of the field and of the challenges and goals that lay ahead. \smallskip The Report is organised as follows: we start in Section \ref{sec:bbp} with a brief overview of the basic features of the inspiral, merger and post-merger of binary systems of neutron stars. This is then followed in Section \ref{sec:ms} by a succinct reminder of the most common formulation of the set of equations needed to simulate the dynamics of BNSs, while the problem of computing initial data is reviewed in Section \ref{sec:ID}. With Section \ref{sec:ph} we will start our review of the progress in simulations in pure hydrodynamics, leaving treatments that include magnetic fields and neutrino transport to Section \ref{sec:bph}. There, special attention is given to the ejecta, which are thought to produce heavy elements and electromagnetic emission in terms of a macronova signal. Finally, Section \ref{sec:atas} is dedicated to the discussion of more advanced techniques and scenarios, which include: high-order numerical methods, the dynamics of BNSs in alternative theories of gravity, as well as the dynamics of binary neutron stars in relativistic collisions. A concluding Section \ref{sec:sao} will summarise the status of research and its future prospects. \medskip We here use a spacelike signature $(-,+,+,+)$ and a system of units in which $c=G=M_\odot=1$ (unless explicitly shown otherwise for convenience). Greek indices are taken to run from $0$ to $3$, Latin indices from $1$ to $3$ and we adopt the standard convention for the summation over repeated indices. \lrn{Finally, reported below is also a quick list of the acronyms adopted in the paper: \begin{tabbing} \hglue 0.15truecm \= \hglue 2.1truecm \= hglue 2.0truecm \kill \> {ADM}: \> Arnowitt, Deser, Misner\\ \> {AMR}: \> adaptive mesh refinement\\ \> {BNS}: \> binary neutron stars\\ \> {BSSNOK}: \> Baumgarte, Shapiro, Shibata, Nakamura, Oohara, Kojima\\ \> {CCZ4}: \> conformal and covariant Z4\\ \> {EOB}: \> effective one body\\ \> {EOS}: \> equation of state\\ \> {ET}: \> Einstein Telescope\\ \> {HMNS}: \> hypermassive neutron star\\ \> {HRSC}: \> high resolution shock capturing\\ \> {IMHD}: \> ideal magnetohydrodynamics\\ \> {KHI}: \> Kelvin-Helmholtz instability\\ \> {LIGO}: \> Laser Interferometer Gravitational-Wave Observatory\\ \> {MHD}: \> magnetohydrodynamics\\ \> {MRI}: \> magnetorotational instability\\ \> {PSD}: \> power spectral density\\ \> {RMHD}: \> resistive magnetohydrodynamics\\ \> {SGRB}: \> short gamma-ray burst\\ \> {SMNS}: \> supramassive neutron star\\ \> {SNR}: \> signal-to-noise ratio\\ \> {TOV}: \> Tolman, Oppenheimer, Volkoff\\ \end{tabbing} } \newpage

\label{sec:sao} As anticipated in the Introduction, there is little doubt that this is a particularly exciting and highly dynamical time for research on neutron stars, in general, and on BNS mergers, in particular. In less than 10 years, \ie starting approximately from 2008, a considerable effort by several groups across the world has obtained numerous important results about the dynamics of binary systems of neutron stars, employing a large variety of numerical (in most cases) and analytical (in a few cases) techniques and exploring this process with different degrees of approximation and realism. Altogether, these works have revealed that the merger of a binary system of neutron stars is a marvellous physical laboratory. Indeed, BNS mergers are expected to be behind several fascinating physical processes, which we recall here: {(i)} they are significant sources of gravitational radiation; {(ii)} they could act as possible progenitors for short-gamma-ray bursts (SGRBs); {(iii)} they have the potential to produce electromagnetic and neutrino emission that is visible from enormous distances; {(iv)} they are likely responsible for the production of a good portion (if not all) of the very heavy elements in the Universe. When viewed across this lens, it is quite natural to consider BNS mergers as Einstein's richest laboratory, binding in the same environment highly nonlinear gravitational dynamics with complex microphysical processes and astonishing astrophysical phenomena. The huge progress accomplished over the last ten years has helped trace a broadbrush picture of BNS mergers that has several sound aspects, among which the most robust in our opinion are the following ones:\footnote{In this Section we will intentionally omit references to avoid cluttering the text; all the relevant references can be found in the various Sections covering the topics discussed here.} \begin{itemize} \item Independently of the fine details of the EOS, of the mass ratio or of the presence of magnetic fields, the merger of a binary system of neutron stars eventually leads to a rapidly rotating black hole with dimensionless spin $J/M^2 \simeq 0.7-0.8$ surrounded by a hot accretion torus with mass \lrn{in the range $M_{\rm torus} \sim 0.001 - 0.1\,M_{\odot}$.} Only very low-mass progenitors whose total mass is below the maximum mass of a (nonrotating) neutron star would not produce a black hole. It is unclear whether such progenitors are statistically important. \item The complete gravitational-wave signal from inspiralling and merging BNSs can be computed numerically with precision that is smaller but overall comparable with that available for black holes. \item When considering the inspiral-only part of the gravitational-wave signal, semi-analytical approximations either in the post-Newtonian or EOB approximation, can reproduce the results of numerical-relativity calculations essentially up to the merger. \item The gravitational-wave spectrum is marked by precise frequencies, either during the inspiral or after the merger that exhibit a ``quasi-universal'' behaviour. \lrn{In other words, while the position of the peaks depends on the EOS, it can be easily factored out to obtain EOS-independent relations between the frequencies of the peaks and the properties of the progenitor stars.} \item The result of the merger, \ie the binary-merger product, is a highly massive and differentially rotating neutron star. The lifetime of the binary-merger product depends on a number of factors, including the mass of the progenitors, their mass ratio and EOS, as well as the role played by magnetic fields and neutrino losses. While sufficiently large initial masses can yield a prompt collapse at the merger, smaller masses can lead to a binary-merger product surviving hundreds of seconds and possibly more. \item When considering magnetic fields of realistic strengths endowing the stars prior to the merger, the correction imprinted by them on the gravitational-wave signal during the inspiral are too small to be detected from advanced gravitational-wave detectors. Electromagnetic signals could be produced before the merger, but these are probably too weak to be detected from cosmological distances. \item Magnetic fields are expected to be amplified both at the merger (via Kelvin-Helmholtz instability), after it and before the collapse and after the formation of a black-hole--torus system (in all cases via a magnetorotational instability or a dynamo action converting small-scale fields into large-scale ones). The final and effective amplification of the resulting magnetic fields is still uncertain, although it should be of at least two-three orders of magnitude. \item The interaction of amplified magnetic fields and accretion in the black-hole--torus system leads to the formation of a magnetically confined plasma along the polar directions of the black hole. Under suitable conditions, the plasma in this funnel may be launched to \lrn{ultrarelativistic} speeds (still unobserved in simulations). \item Matter is expected to be ejected both at the merger and subsequently as a result of a combination of processes: tidal and dynamical mass ejection, magnetically driven winds, neutrino-driven winds, shock-heating winds. Overall, the matter ejected from binaries in quasi-circular orbits amounts to $M_{\rm ejected} \sim 0.001 - 0.01\,M_{\odot}$, while binaries in eccentric orbits can yield up to one order of magnitude more. \item The ejected and unbound matter is expected to undergo nuclear transformations that are mediated by the emission and absorption of neutrinos. Rapid neutron-capture processes ($r$-processes) will then lead to nucleosynthetic yields that are insensitive to input physics or merger type in the regions of the second and third $r$-process peaks, matching the Solar abundances surprisingly well. However, first-peak elements are difficult to explain without invoking contributions from either neutrino and viscously-driven winds operating on longer timescales after the merger, or from core-collapse supernovae. \item The radioactive decay of the ejected matter or its interaction with the interstellar medium are likely to yield afterglows in the infrared or radio bands that are expected to follow the merger after timescales that go from several days to years. \end{itemize} Note that many of the aspects listed above are robust but have been addressed mostly at a rather qualitative level, with precisions that range from ``a-factor-of-a-few'' up to ``order-of-magnitude'' estimates. Furthermore, these results can be seen as the low-hanging fruits of a tree that still has a number of results to offer, although these will require an equal, if not larger, investments of effort, microphysical and numerical developments, and, of course, of computer time. Among the most pressing and exciting open issues we certainly list the following ones: \begin{itemize} \item Consistent initial data for magnetised and arbitrarily spinning neutron-star binaries. \item Semi-analytical and faithful description of the complete gravitational-wave signal, from the inspiral to the formation of a black-hole or stable neutron star. \item Robust and accurate estimate of the critical mass to prompt collapse and of the survival time of the binary-merger product. \item Robust and accurate estimate of the processes mediating the accretion of the torus formed around the black hole, hence obtaining reliable measurements of the timescale for accretion. \item Assessment of the role that turbulence and instabilities play in the amplification of the progenitor magnetic field and determination of the final strengths to be expected. \item Robust and accurate sub-grid modelling for a realistic simulation of the magnetic-field dynamics at the smallest scales. \item Assessment of a possible dynamo action occurring in a long-lived binary-merger product and leading to the generation of an ordered and large-scale magnetic field. \item Determination of the microphysical processes leading to the formation and launching of an ultra-relativistic jet. \item Determination of the acceleration sites of charged particles to ultrarelativistic energies and calculation of the energy distribution functions. \item Determination of the role played by neutrino losses in launching the jet and in modifying the chemical composition of the ejected material. \item Determination of the relative importance of dynamical tidal torques, magnetic unbalance, neutrino emission, or shock heating, for the ejection of matter from the system. \item Robust and accurate determination of the physical and chemical properties of material ejected from the whole merger process. \item Quantitative and accurate predictions of the electromagnetic signal produced by the merger, either directly or indirectly as afterglows. \end{itemize} In conclusion, if the first direct detection of the gravitational-wave signals from binary systems of black holes has officially given birth to the era of gravitational-wave astronomy and has, once again, emphasised general relativity as the best theory of gravitation known, the huge advances that are expected to come in the next few years on the physics and astrophysics of BNSs will help lift many of the veils that still cover Einstein's richest laboratory.

1607.03540

1607

1607.03149.txt

We analyse the Baryon Acoustic Oscillation (BAO) signal of the final Baryon Oscillation Spectroscopic Survey (BOSS) data release (DR12). Our analysis is performed in Fourier-space, using the power spectrum monopole and quadrupole. The dataset includes $1\,198\,006$ galaxies over the redshift range $0.2 < z < 0.75$. We divide this dataset into three (overlapping) redshift bins with the effective redshifts $\zeff = 0.38$, $0.51$ and $0.61$. We demonstrate the reliability of our analysis pipeline using N-body simulations as well as $\sim 1000$ MultiDark-Patchy mock catalogues, which mimic the BOSS-DR12 target selection. We apply density field reconstruction to enhance the BAO signal-to-noise ratio. By including the power spectrum quadrupole we can separate the line-of-sight and angular modes, which allows us to constrain the angular diameter distance $D_A(z)$ and the Hubble parameter $H(z)$ separately. We obtain two independent $1.6\%$ and $1.5\%$ constraints on $D_A(z)$ and $2.9\%$ and $2.3\%$ constraints on $H(z)$ for the low ($\zeff=0.38$) and high ($\zeff=0.61$) redshift bin, respectively. We obtain two independent $1\%$ and $0.9\%$ constraints on the angular averaged distance $D_V(z)$, when ignoring the Alcock-Paczynski effect. The detection significance of the BAO signal is of the order of $8\sigma$ (post-reconstruction) for each of the three redshift bins. Our results are in good agreement with the Planck prediction within $\Lambda$CDM. This paper is part of a set that analyses the final galaxy clustering dataset from BOSS. The measurements and likelihoods presented here are combined with others in~\citet{Alam2016} to produce the final cosmological constraints from BOSS.

\label{sec:intro} The baryon acoustic oscillation (BAO) signal in the distribution of galaxies is an imprint of primordial sound waves that have propagated in the very early Universe through the plasma of tightly coupled photons and baryons (e.g.~\citealt{Peebles:1970ag, Sunyaev:1970eu}). The corresponding BAO signal in photons has been observed in the Cosmic Microwave Background (CMB) and has revolutionised cosmology in the last two decades~\citep[e.g.,][]{Ade:2015xua}. The BAO signal has a characteristic physical scale that represents the distance that the sound waves have traveled before the epoch of decoupling. In the distribution of galaxies, the BAO scale is measured in angular and redshift coordinates, and this observational metric is related to the physical coordinates through the angular diameter distances and Hubble parameters, which in turn depend on the expansion history of the Universe. Therefore, comparing the BAO scale measured in the distribution of galaxies with the true physical BAO scale, i.e., the sound horizon scale that is independently measured in the CMB, allows us to make cosmological distance measurements to the effective redshift of the distribution of galaxies. With this ``standard ruler'' technique one can map the expansion history of the Universe~\citep[e.g.,][]{Hu:1996,Eisen:2003,Blake:2003rh,Linder:2005in,Hu:2003,Seo:2003pu}. While the BAO feature itself can be isolated from the broadband shape of any galaxy clustering statistic quite easily due to its distinct signature, it is still subject to several observational and evolutionary non-linear effects, which damp and shift the BAO feature, thereby biasing such a measurement if ignored~\citep[e.g.,][]{Meiksin:1999,Seo:2005,Crocce:2007dt,Seo:2009,Matsubara:2008wx,Mehta:2011xf,Taruya:2009ir}. In redshift space, the signal-to-noise ratio of the power spectrum is boosted along the line-of-sight due to the linear Kaiser factor $(1+\beta\mu^2)^2$~\citep{Kaiser:1987qv}, but also suffers the non-linear redshift-space distortion effects, which cause additional smearing of the BAO feature along the line of sight. The BAO method has been significantly strengthened by~\citet{Eisenstein:2006nk}, who showed that non-linear degradation effects are reversible by undoing the displacements of galaxies due to bulk flow that are the very cause of the structure growth and redshift-space distortions. This density field reconstruction technique has been tested against simulations and adopted in current galaxy survey data analyses (e.g.~\citealt{Padmanabhan:2012hf}). In this paper, we will apply this technique. The galaxy BAO signal was first detected in the SDSS-LRG~\citep{Eisenstein:2005su} and 2dFGRS~\citep{Percival:2001hw, Cole:2005sx} samples. The WiggleZ survey extended these early detections to higher redshifts~\citep{Blake:2011en,Kazin:2014qga}, while the 6dFGS survey measured the BAO signal at $z=0.1$~\citep{Beutler:2011hx}. Recently the BAO detection in the SDSS main sample at $z = 0.15$ was reported in~\cite{Ross:2014qpa}. The first analysis of the Baryon Oscillation Spectroscopic Survey (BOSS) dataset in DR9~\citep{Anderson:2012sa} presented a $1.7\%$ constraint on the angular averaged distance to $z=0.57$, which has been improved to $1\%$ with DR11~\citep{Anderson:2013zyy}. The LOWZ sample of BOSS has been used in~\citet{Tojeiro:2014eea} to obtain a $2\%$ distance constraint. While the BAO technique has now been established as a standard tool for cosmology, the anisotropic Fourier-space analysis has been difficult to implement because of the treatment of the window function. To simplify the window function treatment it often has been assumed that the window function is isotropic, which simplifies its treatment considerably. However, the window functions of most galaxy surveys are anisotropic and this can introduce anisotropies by re-distributing power between the multipoles and potentially bias cosmological measurements. The first self-consistent Fourier-space analysis which does not put such assumptions on the window function was presented in~\citet{Beutler:2013yhm} using BOSS-DR11, which focused on constraining redshift-space distortions and the Alcock-Paczynski effect. Here we follow this earlier analysis with a few modifications and present a BAO-only analysis, marginalising over the broadband power spectrum shape. Our companion paper~\citep{Beutleretal2} (from now on B16) goes beyond BAO, studying the additional cosmological information of redshift-space distortions. This paper uses the combined data from BOSS-LOWZ and BOSS-CMASS, covering the redshift range from $z=0.2$ to $z=0.75$. The dataset also includes additional data from the so called 'early regions', which have not been included before (see~\citealt{Alam2016} for details). BAO measurements obtained using the monopole and quadrupole correlation function are presented in~\citet{Ross2016}, while~\citet{Vargas-Magana2016} diagnoses the level of theoretical systematic uncertainty in the BOSS BAO measurements. Measurements of the rate of structure growth from the RSD signal are presented in~\citet{Beutleretal2},~\citet{Grieb:2016},~\citet{Sanchez2016} and~\citet{Satpathy2016}. \citet{Alam2016} combines the results of these seven papers (including this work) into a single likelihood that can be used to test cosmological models. The paper is organised as follows. In \S~\ref{sec:data}, we introduce the BOSS DR12 dataset. In \S~\ref{sec:estimator}, we present our anisotropic power spectrum estimator, followed by a description of our window function treatment in \S~\ref{sec:win}. In \S~\ref{sec:mocks} we present the MultiDark-Patchy mock catalogues, which are used to obtain a covariance matrix, and in \S~\ref{sec:model} we introduce our power spectrum model. In \S~\ref{sec:sys} we test our power spectrum model using N-body simulations and the MultiDark-Patchy mock catalogues. In \S~\ref{sec:analysis} we present the data analysis, followed by a discussion of the results in \S~\ref{sec:dis}. We conclude in \S~\ref{sec:conclusion}. The fiducial cosmological parameters which are used to convert the observed angles and redshifts into co-moving coordinates and to generate linear power spectrum models as input for the power spectrum templates, follow a flat $\Lambda$CDM model with $\Omega_m=0.31$, $\Omega_bh^2=0.022$, $h=0.676$, $\sigma_8=0.824$, $n_s=0.96$, $\sum m_{\nu} = 0.06\,$eV and $r_s^{\rm fid} = 147.78\,$Mpc.

\label{sec:conclusion} We have measured the power spectrum multipoles from the final BOSS DR12 dataset in three (overlapping) redshift bins, covering the total redshift range $0.2 < z < 0.75$. Our analysis focuses on measuring the isotropic and anisotropic Baryon Acoustic Oscillation signal in Fourier-space. Our main results are: \begin{enumerate} \item We measure the power spectrum monopole and quadrupole, accounting for window function, aliasing and discreteness effects and extract the BAO information by marginalising over the broadband shape of the power spectrum. We validate our analysis pipeline using two sets of N-body simulations as well as the MultiDark-Patchy mock catalogues. \item Fitting the monopole and quadrupole between $k = 0.01$ - $0.30\hMpc$ produces a constraint on the Hubble parameter of $H(z)r_s/r_s^{\rm fid} = 79.3\pm2.8\,$km\,s$^{-1}$Mpc$^{-1}$ and a constraint on the angular diameter distance of $D_A(z)r_s^{\rm fid} = 1088\pm23\,$Mpc for the low-redshift bin and $H(z)r_s/r_s^{\rm fid} = 98.9\pm2.3\,$km\,s$^{-1}$Mpc$^{-1}$ and $D_A(z)r_s^{\rm fid} = 1433\pm21\,$Mpc for the high-redshift bin (see Table~\ref{tab:results} for a complete summary of the results). While the high-redshift bin is in good agreement with previous results from the CMASS sample, our low-redshift constraint is significantly improved compared to previous studies. Our results are included in~\citet{Alam2016}, where a detailed study of the cosmological implications is performed. \item Ignoring the Alcock-Paczynski effect we can constrain the angular averaged distance $D_V$, for which we obtain a $1\%$ and a $0.88\%$ constraints at the effective redshifts of $z_{\rm eff} = 0.38$ and $0.61$, respectively. \item The detection significances of the BAO signal are $3.4$, $4.2$ and $4.2\sigma$ before applying density field reconstruction for the low, middle and high redshift bins, respectively, and increases to $7.9$, $8.0$ and $8.2\sigma$ after density field reconstruction. \end{enumerate} \citet{Alam2016} combines our measurements with the corresponding correlation function measurements of~\citet{Ross2016} and the growth of structure measurements of~\citet{Beutleretal2},~\citet{Grieb:2016},~\citet{Sanchez2016} and~\citet{Satpathy2016} into a final BOSS likelihood and investigates the cosmological implications. %Mention the public url here

1607.03149

1607

1607.01290_arXiv.txt

The majority of massive disk galaxies, including our own, have stellar bars with vertically thick inner regions -- so-called ``boxy/peanut-shaped'' (B/P) bulges. The most commonly suggested mechanism for the formation of B/P bulges is a violent vertical ``buckling'' instability in the bar, something that has been seen in $N$-body simulations for over twenty years, but never identified in real galaxies. Here, we present the first direct observational evidence for ongoing buckling in two nearby galaxies (NGC~3227 and NGC~4569), including characteristic asymmetric isophotes and (in NGC~4569) stellar-kinematic asymmetries that match buckling in simulations. This confirms that the buckling instability takes place and produces B/P bulges in real galaxies. A toy model of bar evolution yields a local fraction of buckling bars consistent with observations if the buckling phase lasts $\sim$ 0.5--1 Gyr, in agreement with simulations.

% Approximately 60--70\% of disk galaxies in the local universe have stellar bars \citep[e.g.,][]{eskridge00,menendez-delmestre07}. A wide variety of observational evidence indicates that many bars are vertically thickened in their inner regions, appearing as ``boxy'' or ``peanut-shaped'' (B/P) bulges when seen edge-on; this includes our own Galaxy, whose bulge is mostly if not entirely part of its bar \citep[e.g.,][]{shen10,di-matteo14}. Edge-on galaxies with B/P bulges show gas and stellar kinematics consistent with a rotating bar in the disk plane \citep{kuijken95,bureau99a,merrifield99,veilleux99, chung-bureau04}; moderately inclined barred galaxies show isophotes consistent with the projection of B/P bulges within the bars \citep{bettoni94,quillen97,athanassoula06,erwin-debattista13}; and face-on barred galaxies show kinematic and morphological signatures of B/P bulges as well \citep{mendez-abreu08,laurikainen14}. Recent studies suggest that failing to account for the presence of B/P bulges can lead to significantly overestimating the luminosities and masses of ``classical'' (spheroidal) bulges in disk galaxies \citep{laurikainen14,athanassoula15}. This can potentially bias our understanding of how bulges are related to other galaxy properties, including the key correlations between supermassive black holes and bulges \citep[e.g.,][]{kormendy13}. Understanding the formation of B/P bulges is thus an important part of understanding and constraining models of galaxy (and black hole) evolution. The most frequently invoked mechanism for forming these structures is the buckling instability of the bar, a brief but violent vertical instability which occurs (in simulations) not long after the bar forms. In simulations, the formation of the bar increases the radial velocity dispersion of stars in the disk; this leads to a highly anisotropic velocity dispersion tensor and the vertical destabilization of the bar \citep{raha91,merritt94}. Following a phase of asymmetric vertical buckling, the inner part of the bar settles into the more vertically symmetric form of a B/P bulge. Despite over twenty years of simulations which show buckling, no observed buckling has (yet) been reported for real galaxies, a situation called ``puzzling'' in the review of \citet{shen16}. An alternative model proposes that B/P bulges form by the trapping of single orbits into vertical resonances, leading to more gradual, vertically \textit{symmetric} growth \citep{combes90,quillen02,debattista06,berentzen07,quillen14}. In some simulations, the presence of gas weakens or prevents buckling, while still allowing symmetric bar thickening \citep{berentzen98,debattista06,berentzen07,wozniak09}. Thus, it is not clear that real galaxies must suffer the buckling instability. There are also no clear, strong differences due to the different formation mechanisms in the resulting end-stage B/P bulges, making it difficult to determine from observations of existing B/P bulges how they were formed. In this Letter, we present evidence for ongoing buckling in the bars of two local spiral galaxies (NGC~3227 and NGC~4569), thus demonstrating that buckling of bars definitely occurs in real galaxies. We also argue that the observed fraction of buckling bars at $z = 0$ is at least broadly consistent with most (or even all) B/P bulges being the result of buckling, if the buckling phase lasts $\sim 0.5$--1 Gyr -- as is predicted by $N$-body simulations. \begin{figure*} \centering \hspace*{-0.5cm}\includegraphics[scale=1.02]{Figure_sims_pre-during-and-post} \caption[]{ Edge-on and inclined views of $N$-body simulations before, during, and after vertical buckling. Panels a--c show log-scaled isodensity contours of Simulation~C for edge-on (upper sub-panels, with bar perpendicular to line of sight) and inclined views (lower sub-panels, $i = 60\arcdeg$, bar oriented 30\arcdeg{} from line of nodes before inclining galaxy). \textbf{a.} Before buckling, showing the symmetric, vertically thin bar. \textbf{b.} During buckling: vertical asymmetry (upper sub-panel) translates to asymmetric, trapezoidal inner region in lower sub-panel (red contours) and outer-bar ``spurs'' (green), offset in same direction (arrows) from major axis of inner region (cyan line) in lower sub-panel. \textbf{c.} After buckling: the symmetric boxy/peanut-shaped (B/P) bulge projects to rectangular inner contours (red) and counter-offset spurs (green contours, arrows); the projected bar now has 180\arcdeg{} rotational symmetry about the galaxy center. The small inset panels show cartoon versions of the basic buckling and post-buckling projected morphologies (red trapezoid/box + green spurs). \textbf{d--f.} As for lower sub-panel of \textbf{b}, but now showing Simulation~C later in the buckling process (\textbf{d}) and Simulations~B and D during their buckling phases (\textbf{e, f}); all three are seen with $i = 65\arcdeg$. }\label{fig:sims-pre-during-and-post} \end{figure*} \begin{figure*} \centering \hspace*{-1.2cm}\includegraphics[scale=1.0]{Figure_galaxies-vs-sims_6panel} \caption[]{$N$-body simulations during and after buckling, along with real galaxies seen at similar orientations. \textbf{a.} Simulation A after buckling: the isodensity contours show the characteristic box + offset spurs with 180\arcdeg{} rotational symmetry. \textbf{b.} $H$-band isophotes of symmetric-B/P galaxy NGC~3185 \citep{erwin-debattista13}. \textbf{c.} \textit{Spitzer} 3.6~\micron{} isophotes of symmetric-B/P galaxy NGC~3627 \citep{kennicutt03}. \textbf{d.} Simulation C \textit{during} buckling. \textbf{e.} \textit{Spitzer} 3.6~\micron{} isophotes of NGC~4569 \citep{kennicutt03}. \textbf{f.} $K$-band isophotes of NGC~3227 \citep{mrk97}. All images are rotated to place disk line of nodes horizontal; all isophotes are logarithmically scaled. Red contours outline approximate trapezoid (buckling) or boxy (post-buckling) B/P regions, green contours outline outer (vertically thin) bar spurs, and dashed cyan lines show trapezoid/box major axes. \label{fig:n3227-n4569-vs-sim}} \end{figure*}

1607.01290

1607

1607.06365_arXiv.txt

There is strong observational evidence indicating a time lag of order of some 100\,Myr between the onset of starburst and AGN activity in galaxies. Dynamical time lags have been invoked to explain this. We extend this approach by introducing a viscous time lag the gas additionally needs to flow through the AGN's accretion disc before it reaches the central black hole. Our calculations reproduce the observed time lags and are in accordance with the observed correlation between black hole mass and stellar velocity dispersion.

Motivated by, for instance, observed correlations between the mass of an AGN's central black hole and the host galaxy's velocity dispersion \citep[e.g.,][]{2000_Gebhardt_Bender_Bower} and between black hole mass and bulge mass \citep[e.g.,][]{1995ARA&A..33..581K}, there is an ongoing debate whether, and if so, how starbursts and AGN are connected to each other. \citet{2005_Di-Matteo_Springel_Hernquist}, for instance, explain such correlations as due to a thermal AGN feedback that heats the gas of the galaxy and thus prevents further star formation and AGN activity: More massive galaxies have a deeper gravitational potential well, thus the black hole has to gain more mass before its luminosity is capable of expelling the gas from the galaxy and quenching star formation and AGN activity. This then leads to the velocity dispersion and the bulge mass, resp., to be related to the black hole mass. In these simulations starburst and AGN activity occur simultaneously, but recent observations show that AGN activity may be delayed with regard to star formation activity by time-scales of 50--250\,Myr \citep[e.g.,][]{2007_Davies_Mueller-Sanchez_Genzel, 2009ApJ...692L..19S, 2010MNRAS.405..933W}. \citet{2012MNRAS.420L...8H} argues that such a time lag can occur for purely dynamical reasons. His high spatial resolution simulations of galaxy mergers show first an inward motion of gas towards the dynamical centre giving rise to (a burst of) star formation. In these models, the gas flowing further inwards can do so only by losing angular momentum by gravitational instabilities. This, in turn, gives rise to a time lag between star formation and AGN activity. We extend this idea by modelling the loss of angular momentum and the ensuing inflow in the framework of an accretion disc scenario. Thus the time lag between starburst and AGN activity consists of a dynamical lag given by the time span the gas needs to reach the accretion disc and a subsequent viscous lag given by the time span the gas needs to flow through the accretion disc until it reaches the black hole. In Section \ref{sec:numerics} we explain our numerical methods and the setup of the merger event that is, in this scenario, responsible for the inflow of gas to the newly forming galactic centre. In Section \ref{sec:results} we present and discuss the general picture that results from our calculations. As our model depends on a number of parameters, we perform a parameter study in Section \ref{sec:param} to show the robustness of our results against parameter changes. In Section \ref{sec:summary} we summarise our findings.

\label{sec:summary} With only one model setup we were able to reproduce three observational findings that have been identified in galaxies: (i) The observed time lag between starburst and AGN activity is, in our work, principally caused by a viscous time lag the gas needs to flow through the AGN's accretion disc until it reaches the central black hole. (ii) Our results match the observed $M_{\rmn{BH}}$-$\sigma$ correlation, but additionally include the aforementioned time lags. As, e.g., \citet{2005_Di-Matteo_Springel_Hernquist} or \citet{2011_Debuhr_Quataert_Ma} have already shown, AGN feedback is responsible for this relation. (iii) The large scatter of the $M_{\rmn{BH}}$-$\sigma$ correlation is, in our work, caused by the continuing evolution of the black hole mass after the merging event. The time lag does not only occur between starburst and AGN activity, but also between the peak in the AGNAR and the peak in the BHAR. Thus it is evident that this time lag is caused by the gas being delayed by the AGN particle, i.e. by the gas needing some time to move through the accretion disc. In addition our parameter study shows that the time lag increases with increasing size of the accretion disc, and decreases with increasing viscosity, indicating that it is indeed a viscous time-scale that delays the gas. Furthermore the parameter study shows that a time lag between starburst and AGN activity of the order of some 100\,Myr can be reproduced for a wide range of parameters and for different resolution levels. Thus our conclusions do not depend on the detailed choice of the parameters, as long as they are taken from within a physically sensible range.

1607.06365

1607

1607.06686_arXiv.txt

In the Solar system, a quasi-satellite is an object that follows a heliocentric path with an orbital period that matches almost exactly with that of a host body (planetary or not). The trajectory is of such nature that, without being gravitationally attached, the value of the angular separation between host and quasi-satellite as seen from the Sun remains confined within relatively narrow limits for time-spans that exceed the length of the host's sidereal orbital period. Here, we show that under these conditions, a quasi-satellite traces an analemma in the sky as observed from the host in a manner similar to that found for geosynchronous orbits. The analemmatic curve (figure-eight-, teardrop-, ellipse-shaped) results from the interplay between the tilt of the rotational axis of the host and the properties of the orbit of the quasi-satellite. The analemma criterion can be applied to identify true quasi-satellite dynamical behaviour using observational or synthetic astrometry and it is tested for several well-documented quasi-satellites. For the particular case of 15810 (1994~JR$_1$), a putative accidental quasi-satellite of dwarf planet Pluto, we show explicitly that this object describes a complex analemmatic curve for several Plutonian sidereal periods, confirming its transient quasi-satellite status.

Objects trapped in a 1:1 mean motion resonance with a host (planetary or not) are classified as co-orbitals of the host, independently of the shape and orientation of their paths (Morais \& Morbidelli 2002); in other words, to be classed as co-orbitals their orbits do not have to resemble that of the host as long as the ratio of their orbital periods equates to almost exactly one. In general, co-orbital configurations are not identified observationally but as a result of the statistical analysis of large sets of numerical integrations. There is, however, a potential exception to this standard approach; a particular type of co-orbital configuration that can be confirmed observationally, the quasi-satellite dynamical state. Here, we study the apparent motion in host-centric equatorial coordinates of known quasi-satellites to show that they trace an analemmatic curve in the sky as observed from the host in a manner similar to that found for geosynchronous orbits. This paper is organized as follows. Section~2 discusses the so-called analemma criterion for quasi-satellites and it includes an extensive exploration of the known quasi-satellite population. The particular case of 15810 (1994~JR$_1$), a putative accidental quasi-satellite of Pluto, is analysed in Section~3 to show that according to the analemma criterion it is a true transient quasi-satellite of Pluto. Results are discussed in Section~4 and conclusions are summarized in Section~5.

In this paper, we have explored a new criterion to identify quasi-satellites. In sharp contrast with the numerical strategies customarily applied in the study of co-orbital bodies, the criterion described here can make direct use of observational astrometric data. Our conclusions can be summarized as follows. \begin{enumerate}[(i)] \item Bona fide quasi-satellites trace paths in the sky which repeat every sidereal period when observed from their hosts. These paths can be described as analemmatic curves similar to those found for geosynchronous orbits. The analemma shifts as the orbit of the quasi-satellite changes over time. \item The existence of this analemmatic behaviour turns quasi-satellites, natural or artificial, into potentially interesting platforms for the future of space exploration. \item The Earth has the largest known number of present-day quasi-satellites, five. Jupiter comes in second place with three. Venus, Saturn, Neptune and dwarf planet Ceres have one each. \item Applying the analemma criterion, Plutino 15810 (1994 JR$_1$) is as good a quasi-satellite as it may get. Therefore, dwarf planet Pluto hosts at least one quasi-satellite at present. \item Asteroid 63252 (2001~BL$_{41}$) is a present-day transient quasi-satellite of Saturn. \item Historically, the first object identified as quasi-satellite (in this case of Jupiter) was an asteroid moving in a comet-like orbit. Unfortunately, this object appears to have been lost since its announcement back in 1973. \end{enumerate}

1607.06686

1607

1607.00775_arXiv.txt

We examine the effect that the magnetic part of the Weyl tensor has on the large-scale expansion of space. This is done within the context of a class of cosmological models that contain regularly arranged discrete masses, rather than a continuous perfect fluid. The natural set of geodesic curves that one should use to consider the cosmological expansion of these models requires the existence of a non-zero magnetic part of the Weyl tensor. We include this object in the evolution equations of these models by performing a Taylor series expansion about a hypersurface where it initially vanishes. At the same cosmological time, measured as a fraction of the age of the universe, we find that the influence of the magnetic part of the Weyl tensor increases as the number of masses in the universe is increased. We also find that the influence of the magnetic part of the Weyl tensor increases with time, relative to the leading-order electric part, so that its contribution to the scale of the universe can reach values of $\sim 1\%$, before the Taylor series approximation starts to break down.

The large-scale expansion of the Universe is usually taken to be dominated by the Newtonian part of the gravitational field. This idea probably originates from the close correspondence between the Friedmann equations of general relativity and the equations that govern Newtonian cosmologies \cite{milne}, but also has strong support from rigourous constructions that are based on perturbative expansions of Einstein's field equations \cite{viraj,virajX}. Nevertheless, there are good reasons to be interested in relativistic effects in cosmology. These include the apparent existence of Dark Energy, as well as the dawn of the age of precision cosmology. Put bluntly, relativistic effects need to be understood in order to have faith in the cosmological models that we use to interpret observational data. Geometrically, the electric part of the Weyl tensor is sufficient to determine the Newtonian part of the free gravitational field. The magnetic part of the Weyl tensor, on the other hand, describes other aspects of the relativistic gravitational field \cite{buchert} (both of these objects are defined in Section \r{sec:formalism}, below). Well known relativistic effects such as frame-dragging and gravitational radiation require a non-vanishing magnetic Weyl tensor in order to exist, but the effects of the magnetic part of the Weyl tensor on the large-scale expansion of the Universe are still largely unknown. This is partly due to an absence of realistic cosmological models that have non-zero magnetic Weyl curvature, and that could therefore be considered to be non-silent \cite{bruni}. This situation is, however, changing. It has recently been demonstrated that cosmological models that contain regularly-arranged discrete masses can generate a non-zero magnetic Weyl tensor, even if no such curvature existed in their initial data \cite{KHB}. The authors of this study considered the effect that this tensor has on the large-scale expansion of a universe that contains eight black holes, using both the leading-order term in a Taylor series expansion, and by numerically integrating the evolution equations of the initially silent space. We extend their study by calculating both the leading-order and next-to-leading-order parts of the relevant series expansion, and by using these results to determine the effect of the magnetic part of the Weyl tensor on the large-scale expansion of all cosmological models that contain regularly arranged discrete masses in a closed cosmology. The phenomenology of such models is interesting, as they can be considered a first approximation to the type of universe within which we actually live. More mathematically, they provide a tractable way to formulate $n$-body cosmology as a relativistic initial value problem. They therefore provide an ideal arena within which to study relativistic gravitational effects in cosmology. Of course, the magnetic part of the Weyl tensor is a frame-dependent object, and its magnitude will change when different sets of observers are considered. However, in cosmology it is natural to consider sets of time-like curves that are both geodesic, and (in some sense) comoving with the objects that exist within the space-time. It is with respect to observers following such a set of curves that one can talk about ``cosmological expansion'', and it is with respect to the same set of curves that we will talk about ``the magnetic part of the Weyl tensor''. The influence of the magnetic part of the Weyl tensor on the cosmological expansion is of special interest because virtually all cosmological solutions of Einstein's equations are derived under assumptions that force it to be zero (for cosmologically interesting congruences of curves). Such situations are, however, highly unrealistic, as observers that follow the geodesic set of time-like curves that describe the world-lines of real galaxies will certainly experience the consequences of this part of the curvature of space-time, which is in general non-zero. We must therefore precisely quantify the effects of the magnetic part of the Weyl tensor on the large-scale expansion of space, if we are to fully understand the recessional velocity of galaxies in the real Universe (and all of its associated consequences). In Section \r{sec:models} we describe in detail the cosmological models that we will be using in this study. They consist of regularly-arranged discrete masses in a closed space, and are formulated as a relativistic initial value problem. In Section \r{sec:formalism} we then introduce the compact formalism that will be used in the rest of the paper. This starts by using a 1+3 decomposition of the geometric and kinematic quantities involved, and finishes with a 1+1+2 decomposition that can be efficiently used to exploit the symmetries of the problem. The explicit form of the symmetry restrictions are then determined in Section \r{sec:reflection}, where it is shown that the influence of the magnetic part of the Weyl tensor on the relevant evolution equations does not necessarily vanish. Section \r{calcs} then contains some lengthy calculations to determine the coefficients of a Taylor series expansion that can be used to include the effect of the magnetic Weyl tensor on the large-scale expansion of space. In Section \r{sec:results} we present the numerical results for each of the lattice universes that we will be considering, before concluding in Section \r{sec:discussion}. Throughout the paper we use the first half of the Latin alphabet to denote space-time coordinate indices, and the second half to denote spatial coordinate indices. Greek letters are reserved to denote spatial frame indices, where they are required.

\l{sec:discussion} We have considered the effect of $H_{ab}$ on the evolution of LRS curves in lattice models of the universe. These models treat the matter in the Universe as a collection of point-like sources, and allow the formulation of cosmology as an initial value problem. They are therefore particularly well suited to the study of relativistic effects in cosmology, including the study of the evolution and effects of the magnetic part of the Weyl tensor. We find that although the initial data of our models is silent (with $H_{ab}=0$), the evolution of the space is not silent. In particular, a $\curl H_{ab}$ term appears in the evolution equation for the electric part of the Weyl tensor, which in turn acts as the source for the evolution of LRS curves. This result was identified and studied numerically by Korzy\'{n}ski, Hinder and Bentivegna for the particular case of a universe that contains eight black holes \cite{KHB}. We extend their study by calculating the leading-order and next-to-leading order terms in a Taylor series approximation that can be used to incorporate the effects of $H_{ab}$ on LRS curves, and by applying it to all possible regular arrangements of black holes in a closed cosmological model. The inclusion of the next-to-leading order term, in our study, allows us to estimate when the series expansion stops converging, and when numerical techniques need to be employed instead. We find that the effect of $H_{ab}$ is small, while the series expansion remains valid, but grows with time. In particular, we find that while the effect of the $H_{ab}$ on the expansion of LRS curves can be at the level of $1\%$ when the number of masses in the Universe is small (as in \cite{KHB}), but that this number decreases as the number of masses in the universe is increased (when comparing at the same time, measured in terms of $m$). While the effect on the expansion rate is small for larger lattices, however, it does appear that the effect of $H_{ab}$ on the expansion is, in some sense, cumulative. That is, the difference in the curve length that results from including the effect of $H_{ab}$ tends to increase over time, as the magnitude of the magnetic Weyl tensor increases over time. When comparing lattices at the same cosmological time ({\it i.e.} at the same time, as measured in units of $n \times m$), it therefore appears that the effect of $\curl H_{ab}$ on the evolution of the LRS curve increases as the number of masses in the universe is increased. Whether or not similar result holds in cosmological models that expand for all time, rather than re-collapsing, remains to be seen. It should be noted, however, that while the effects we find are always small (less than $1\%$, in most cases), the Taylor series approximation used to derive them usually breaks down on time scales that are shorter than the age of the universe. This is especially true in the larger lattices, where it can be seen that increasing the number of masses in the universe results in the series approximation breaking down at earlier cosmological times. This result holds for all of the LRS curves that we studied, but seems to be especially true of the cell edges (which, unfortunately, are probably the best indicators of the scale of the cosmology as a whole). To reliably follow the evolution of these curves any further will require more advanced techniques. It is interesting to see that the effect of higher-order terms in Einstein's equations can lead to non-negligible effects in the large-scale expansion of space. One may note, for example, that it is not until we get to order $t^6$ in the Taylor series expansion in Eq.\ (\ref{da1}) that the influence of $\curl H_{ab}$ becomes non-zero at all. In a weak-field expansion of the gravitational field, about a suitably chosen background, this would probably correspond to quite a high order in perturbations. Nevertheless, the terms involved grow rapidly over cosmological time-scales, until they come close in magnitude to the leading-order terms represented by $E_{ab}$. This is remiscent of the gravitational wave memory effect \cite{memory}, where the small effects of non-linear gravity accumulate over time until (in the case of gravitational waves from astrophysical sources) they are comparable with the magnitude of the linear, leading-order terms.

1607.00775

1607

1607.05961_arXiv.txt

We investigate the high-speed ($v >$ 1000~km~s$^{-1}$) extreme-ultraviolet (EUV) wave associated with an X1.2 flare and coronal mass ejection (CME) from NOAA active region 11283 on 2011 September 6 (SOL2011-09-06T22:12). This EUV wave features peculiar on-disk signatures, in particular we observe an intermittent ``disappearance'' of the front for 120~s in SDO/AIA 171, 193, 211~{\AA} data, whereas the 335~\AA~filter, sensitive to hotter plasmas (T$\sim$ 2.5~MK), shows a continuous evolution of the wave front. The eruption was also accompanied by localized coronal dimming regions. We exploit the multi-point quadrature position of SDO and STEREO-A, to make a thorough analysis of the EUV wave evolution, with respect to its kinematics and amplitude evolution and reconstruct the SDO line-of-sight (LOS) direction of the identified coronal dimming regions in STEREO-A. We show that the observed intensities of the dimming regions in SDO/AIA depend on the structures that are lying along their LOS and are the combination of their individual intensities, e.g. the expanding CME body, the enhanced EUV wave and CME front. In this context, we conclude that the intermittent disappearance of the EUV wave in the AIA 171, 193, 211~\AA~filters, which are channels sensitive to plasma with temperatures below $\sim$ 2~MK is also caused by such LOS integration effects. These observations clearly demonstrate that single-view image data provide us with limited insight to correctly interpret coronal features.

Coronal mass ejections (CMEs) are large-scale coronal magnetic field structures expelled into the heliosphere. They are the main causes of space weather disturbances since they may drive interplanetary shocks and produce geomagnetic storms. Accompanied with CMEs we observe activities low in the solar atmosphere including filament eruptions, flares, large-scale coronal EUV waves, and coronal dimmings. EUV waves are large-scale disturbances propagating through the solar atmosphere, observed as moving fronts of increased coronal EUV emission. They propagate at typical speeds of 200--400\,km\,s$^{-1}$ \citep[][]{Klassen:2000, Thompson:2009, Muhr:2014} but on rare occasions, EUV wave events with speeds in excess of 1000\,km\,s$^{-1}$ have been reported \citep[][]{Nitta:2013}. Several models have been proposed to explain the nature of EUV waves, falling into three categories: wave, non-wave and hybrid models. Most wave models interpret the propagating signatures as fast-mode magnetosonic waves \citep[][]{Thompson:1999, Wang:2000, Wu:2001, Warmuth:2001, Ofman:2002}, while in the non-wave models they are explained as disturbances due to sucessive restructuring of magnetic field lines during the erupting CME \citep[][]{Delannee:1999, Chen:2002,Attrill:2007}. Hybrid models include both interpretations resulting in a bimodal picture, where a fast-mode wave travels ahead of a slower inner component due to magnetic field reconfiguration caused by the erupting CME \citep[][]{Zhukov:2004}. A detailed discussion on the various theories and the evidence for and against them may be found in recent reviews \citep[e.g.][]{Warmuth:2010, Gallagher:2011, Zhukov:2011, Patsourakos:2012, Liu:2014, Warmuth:2015}. The relationship between CMEs and EUV waves was investigated by several statistical studies \citep[][]{Biesecker:2002, Cliver:2005, Nitta:2013, Nitta:2014, Muhr:2014}. It was found that fast and wide CMEs are in general accompanied by well-observed EUV waves that are associated with shocks, i.e. related to type II radio bursts \citep[][]{Cliver:2005}. The unprecedented multi vantage-point observations from the \textit{Solar TErrestrial RElations Observatory} \citep[STEREO,][]{Kaiser:2008} and the high-cadence imagery by the \textit{Atmospheric Imaging Assembly} \citep[AIA,][]{Lemen:2012} on-board the \textit{Solar Dynamics Observatory} \citep[SDO,][]{Pesnell:2012} improved our understanding of how the EUV wave formation and kinematics are related to the CME dynamics. Initially the EUV wave front is found to be closely attached to the laterally expanding CME flanks. After an impulsive driving phase, the CME flanks decelerate, and the EUV wave subsequently becomes freely propagating with a velocity close to the fast magnetosonic speed in the quiet corona \citep[][]{Warmuth:2004, Long:2008, Veronig:2008, Patsourakos:2009, Patsourakos:2009a, Kienreich:2009, Veronig:2010, Olmedo:2012, Cheng:2012}. Quadrature observations of the two STEREO spacecraft unambiguously revealed that the EUV wave is propagating ahead of and is driven by the laterally expanding CME flanks \citep[][]{Patsourakos:2009a, Kienreich:2009, Ma:2011}. These findings were also confirmed by MHD simulations of CME eruptions and associated EUV waves \citep[][]{Pomoell:2008, Cohen:2009, Downs:2011, Downs:2012}. Coronal dimmings are regions of reduced emission in the low corona observed in EUV \citep[][]{Thompson:1998,Thompson:2000} and soft X-rays (SXR; \citealt{Hudson:1996,Sterling:1997}). The appearance of these dimming regions is in general interpreted as density depletion caused by the evacuation of plasma during the CME expansion \citep[][]{Hudson:1996,Thompson:1998, Harrison:2000}. This interpretation is supported by the simulatenous and co-spatial observations of coronal dimmings in different wavelengths \citep[e.g.][]{Zarro:1999}, spectroscopic observations showing plasma outflows in dimming regions \citep[][]{Harra:2001,Attrill:2010,Tian:2012} and studies on dimming/CME mass relations \citep[][]{Sterling:1997, Wang:2002, Harrison:2003, Zhukov:2004, Mandrini:2007, Aschwanden:2009, Miklenic:2011}. In the literature, two different types of dimmings are defined, core (or twin) dimmings and secondary (or remote) dimmings, respectively. Core dimmings are stationary, localized regions that are often present on both sides of an erupting configuration, in opposite magnetic polarity regions. In relatively simple cases they are interpreted to mark the footpoints of the ejected fluxrope \citep[][]{Sterling:1997, Thompson:2000, Webb:2000, Mandrini:2005,Mandrini:2007, Temmer:2011}. The more shallow secondary dimmings are diffuse and can extend to significant distances from the source region. They are often observed to follow behind a propagating EUV wavefront \citep[][]{Delannee:1999, Wills-Davey:1999, Thompson:2000, Attrill:2007, Mandrini:2007,Muhr:2011}. Therefore, they could be rarefaction regions that develop behind a compressive wave \citep[e.g.][]{Wu:2001, Muhr:2011, Downs:2012} or formed due to the plasma evacuation behind the flux rope and overlying fields that are erupting. In this paper, we analyze the flare-CME event and associated large-scale EUV wave that occurred on \mbox{2011 September 6}. Several aspects of this event have been studied before. For instance, \cite{Nitta:2013} investigated the EUV wave kinematics using high-cadence observations from SDO/AIA, and determined this event to be a high-speed EUV wave with a velocity of $v\approx$\,1250\,km\,s$^{-1}$. Only 6 out of 171 events reported by \cite{Nitta:2013} were identified in this high-speed regime ($v>$1200\,km\,s$^{-1}$). The globally propagating EUV wave caused oscillations of filaments and the launch of a jet \citep{Shen:2014}. These results indicate that EUV waves may serve as agents for linking successive solar activities. \cite{Jiang:2013} focused on the triggering mechanism of the erupting event using MHD simulations initialized by a non-linear force-free field (NLFFF) extrapolation. \cite{Janvier:2016} studied the morphology and time evolution of photosperic traces of the current density and flare ribbons and compared it with the topological features found by NLFFF modeling. Both, \cite{Jiang:2013} and \cite{Janvier:2016} identified a spine-fan configuration of the overlying field lines, due to the presence of a parasitic positive polarity, embedding an elongated flux rope. The energy accumulation of the flare and the accompanied CME was studied by \cite{Feng:2013}. They found that the calculated free magnetic energy is able to power the flare and the CME, and that both phenomena may consume a similar amount of free energy. \cite{Romano:2015} studied the evolution of the source active region (NOAA 11283) in relation to its recurrent flaring and CME activity. They found that before the occurence of the X2.1 flare-CME event under study here, the shearing motions seem to inject a larger fraction of energy into the corona than the emergence of the magnetic field. In this paper, we investigate this extraordinary EUV wave and associated flare-CME event in detail and fully exploit the quadrature view from SDO and STEREO.

\label{discussion} The EUV wave/CME/coronal dimming event on September 6, 2011 shows several peculiar features. Only by complementing the results of SDO observations with ST-A, we were able to disentangle real signatures from signatures arising from LOS integration and projection effects. Furthermore, we were able to find new aspects for the formation of associated coronal dimming regions. 1) For the first time we identified the ``disappearance" of a segment of the EUV wave front followed by its reappearance and further propagation in the absence of any obvious magnetic structures. In the following we discuss three possible scenarios for this phenomenon: i) Heating at the wave front: \cite{Vanninathan:2015} studied the strong EUV wave associated with the X2.2 flare of 2011 February 15 and showed that the passing EUV front adiabatically compresses the ambient coronal plasma, resulting in an increase in density of about 6-9\% and an associated temperature increase of 5-6\%. Using the method described in \cite{Downs:2011}, we estimated for the event under study the intensity ratio for each AIA channel as a function of temperature, by assuming that the passing EUV wave causes $\sim$8\% increase in density compared to the pre-event corona (consistent with the values derived by \citealt{Veronig:2011, Schrijver:2011} and \citealt{Vanninathan:2015}). This leads to a $\sim$5\% increase in temperature. From the estimated intensity ratios we obtain that for a $\sim$5\% increase in temperature, the intensity in the 171~\AA~and 193~\AA~channels is reducing while that of 211~\AA~and 335~\AA~channels continue to increase. We repeated this test for different densities ranging from 6-12\% increase and obtained the same qualitative results in all cases. This is contradictory to the observed evolution of the amplitude of the wave pulse revealing a decrease in intensity for all four wavelength channels during its ``disappearance'' phase (cf.~Fig.~\ref{fig:amplitude}). This result is inconsistent with the disappearance being caused by heating at the wave front. ii) Wave propagation through an inhomogenous medium: EUV waves are known to be affected by inhomogeneities in the corona (in terms of Alfv\'en velocity), such as active regions, coronal holes or filaments. Encountering a region of high Alfv\'en speed would alter the amplitude of the wave pulse. To test this hypothesis, we checked for magnetic obstacles in SDO/HMI LOS magnetograms along the propagation path, which would correspond to regions of enhanced Alfv{\'e}n speed. Fig.~\ref{hmi_mean_pos_neg} shows LOS magnetic field and H$\alpha$ data together with the distance range, where the wave front ``disappears" (marked by the red and green fronts). No magnetic flux enhancements or filaments could be identified along the sector of interest. However, a filament is present along the propagation direction of the EUV wave but not at the locations where we observe the wave front disappearance. We note that the wave front passage causes the filament to oscillate, but it does not erupt \citep[][]{Shen:2014}. These observations suggest that the propagation through an inhomogenous medium is not responsible for the disappearance of the wave front. iii) LOS effects: The simultaneous decrease in amplitude for all channels indicates that the decrease in emission results primarily from changes in the density and not changes in temperature. The 171~\AA, 193~\AA, and 211~\AA~filters reveal the highest decrease. These channels are most sensitive to temperatures around 1-2~MK, which corresponds to quiet Sun coronal temperatures. During the CME lift-off such plasma is evacuated, resulting in regions of reduced intensity (coronal dimmings). The observed intensity in SDO/AIA results from the sum of emission along the LOS. Assuming that along a specific LOS, contributions from coronal dimming regions (negative contributions) as well as the EUV wave (positive contributions) are present, the total intensity can be reduced to the background intensity, when summed up. The 335~\AA~channel measures plasma at higher temperatures around $2.5\times10^{6}$~K and is thus less sensitive to the expansion and evacuation of quiet coronal plasma by the erupting CME (cf. Fig.~\ref{fig:event}). \begin{figure} \begin{minipage}{1.0\columnwidth} \centering \includegraphics[width=0.5\textwidth]{./fig10.pdf} \caption{Top: HMI magnetogram. Bottom: GONG H$\alpha$ filtergram showing the location of filaments. Overlaid on both panels is the measured sector (marked in blue), and the wave front determined at the different time steps that span the range of the wave ``disappearance'' (red and green line).} \label{hmi_mean_pos_neg} \end{minipage} \end{figure} 2) Observations from SDO/AIA revealed signatures of different types of coronal dimmings. Using a thresholding technique we were able to distinguish potential core dimming regions from secondary/remote dimming regions. The identified core dimming regions show evacuation of dense plasma, observed as maximum decrease in intensity in the overall identified dimming region from SDO/AIA and \textbf{as the largest decrease} in emission from the ST-A LOS reconstructions. They are located in regions of opposite magnetic polarities and lie close to the eruption site (cf.~Fig.~\ref{dimming}). Therefore, we conclude that these regions point to real footpoints of the erupting fluxrope. For the secondary/remote dimmings (red ROI and lower part of the cyan ellipse in Fig.~\ref{dimming_los}) the reconstruction of the emission from ST-A revealed that whether the observed intensity in SDO/AIA results in a dimming or not, depends strongly on the structures lying along the corresponding LOS. We conclude that projection effects play an important role in the appearance and correct interpretation of coronal dimming regions. 3) The kinematical evolution of the wave is derived to be different when measured from different vantage points. This can be explained by projection effects as shown in Fig.~\ref{fig:stereo}. Obviously, the on-disk wave signatures cover contributions from the outermost bright front of the CME eruption, when the EUV wave is not yet detached from the CME which is driving the wave. \cite{Ma:2009} and \cite{Hoilijoki:2013} found that there exist viewing angles, where the orientation of the LOS is tangential to the erupting dome, producing visible features that can be interpreted as parts of an EUV wave. In a recent study by \cite{Delannee:2014}, using different techniques for deriving the altitude of an EUV wave, the resulting heights span a distance range of 34--154~Mm above the solar limb, which is consistent with our result of $\sim$\,130~Mm for the peak of the wave pulse identified from SDO/AIA. Another important aspect is that the on-disk observations of EUV waves may contain significant intensity contributions of the emission from the expanding CME body integrated along the LOS over several scale heights. The speeds of the EUV wave (as derived from SDO/AIA) and the CME (derived from ST-A with low projection effects) are similar, implying that in both instances we measure the same propagating structure. In addition, the observed ``bending'' of the wave towards the solar East (see movie no 1.\ time range $\sim$~22:24~UT until 22:30~UT), is another feature suggestive of projection effects, as no magnetic obstacles are present at that location, that would cause a reflection or change of propagation direction. Assuming that the ``bending'' results from LOS integration along the outermost front of the CME, we derive that these contributions might come from large heights in the corona. For the time when the wave ``bending'' is clearly observed, the apex of the CME, as derived from ST-A, is at a distance of about~$2-2.5$~R$_{\odot}$. The event studied in this paper is a good example of how the use of single-view image data may limit our ability to correctly interpret coronal features, such as EUV waves and coronal dimmings.

1607.05961

1607

1607.02546_arXiv.txt

Cosmic far-infrared background (CFIRB) is a powerful probe of the history of star formation rate (SFR) and the connection between baryons and dark matter across cosmic time. In this work, we explore to which extent the CFIRB anisotropies can be reproduced by a simple physical framework for galaxy evolution, the gas regulator (bathtub) model. This model is based on continuity equations for gas, stars, and metals, taking into account cosmic gas accretion, star formation, and gas ejection. We model the large-scale galaxy bias and small-scale shot noise self-consistently, and we constrain our model using the CFIRB power spectra measured by {\em Planck}. Because of the simplicity of the physical model, the goodness of fit is limited. We compare our model predictions with the observed correlation between CFIRB and gravitational lensing, bolometric infrared luminosity functions, and submillimetre source counts. The strong clustering of CFIRB indicates a large galaxy bias, which corresponds to haloes of mass $10^{12.5}\Msun$ at $z=2$, higher than the mass associated with the peak of the star formation efficiency. We also find that the far-infrared luminosities of haloes above $10^{12}\Msun$ are higher than the expectation from the SFR observed in ultraviolet and optical surveys.

Cosmic far-infrared background (CFIRB) originates from unresolved dusty star-forming galaxies across cosmic time. In these galaxies, the ultraviolet (UV) photons associated with newly formed, massive stars are absorbed by dust and re-emitted in far-infrared (FIR), and the FIR emission serves as an indicator of the star formation rate (SFR). At the FIR wavelengths ($\sim$100 \micron\ to 1\ mm, also known as submillimetre), most galaxies are unresolved and can only be observed as background intensity fluctuations. These fluctuations contain information about the cosmic star formation history, as well as the dark matter haloes in which the dusty star-forming galaxies are located. Compared with UV, the star formation history from FIR is much less explored because of the limited angular resolutions of the telescopes; thus, CFIRB provides an important piece of the puzzle of the cosmic star formation history. Predicted half a century ago \citep{PartridgePeebles67b,Bond86}, the CFIRB was first discovered by {\em COBE}-FIRAS \citep{Puget96,Fixsen98,Hauser98,Gispert00,HauserDwek01} and subsequently observed by ISO \citep{LagachePuget00,Matsuhara00,Elbaz02}. The anisotropies of CFIRB have been measured by {\em Spitzer} \citep{GrossanSmoot07,Lagache07}, BLAST \citep{Viero09}, SPT \citep{Hall10}, ACT \citep{Hajian12}, {\em Herschel} \citep{Amblard11,Berta11,Viero13}, and {\em Planck} \citep{Planck11CIB, Planck13XXX}. In particular, the angular power spectra of CFIRB provide the luminosity-weighted galaxy bias and thus the information about the mass of the underlying dark matter haloes \citep[e.g.,][]{Viero09,Amblard11,Planck11CIB,DeBernardis12,Shang12,Xia12,Thacker13,Viero13,Planck13XXX}. To date, most of the interpretations of the CFIRB anisotropies are based on phenomenological models with limited physical interpretation. For example, \citet{Addison13} modelled the CFIRB and number counts using general parametrizations for the luminosity function, the spectral energy distribution (SED), and the scale-dependent galaxy bias. On the other hand, \citet{Shang12} implemented a luminosity--mass relation in the halo model to improve the modelling at small scales \citep[also see, e.g.,][]{Viero13}. In addition, \citet{Planck13XXX} provided updated measurements of the CFIRB power spectra as well as new constraints on linear and halo models; however, the SFR density inferred from their halo model appears higher at high redshift when compared with UV and optical observations. In this work, we develop a physical model for the connection between dark matter haloes and dusty star-forming galaxies. We constrain this model using the CFIRB power spectra measured by {\em Planck}. We then compare our model with various FIR/submillimetre galaxy observations. Our model provides a simple, physically-motivated framework to compare and interpret various FIR observations. We apply the gas regulator model, which is based on the continuity equations of gas, stars, and metal \citep[also known as the bathtub or reservoir model, see, e.g.,][]{Bouche10,KrumholzDekel12,Dekel13,Lilly13,DekelMandelker14}, to calculate SFR. We then apply the halo model to calculate the power spectra of CFIRB \citep{ScherrerBertschinger91,Seljak00,CooraySheth02}. We fit the model to the CFIRB anisotropies measured by {\em Planck} \citep{Planck13XXX}. Our model predictions are compared with various IR observations, as well as the cosmic SFR density and cosmic dust mass density constrained by other observations. We find that CFIRB requires high IR luminosity for massive haloes ($\LIR \sim 10^{12}\Lsun$ for haloes of mass above $10^{13} \Msun$); this result is consistent with earlier findings \citep[e.g.,][]{Shang12,Addison13,Bethermin13} but is in excess compared with the SFR constrained by UV and optical. This excess of IR luminosity can be related to heating by old stellar populations. This paper is organized as follows. Section~\ref{sec:bathtub} describes the gas regulator model and provides a quasi-steady-state solution relevant for SFR and dust property. In Section~\ref{sec:halomodel}, we incorporate the gas regulator model into the halo model to calculate observed quantities. In Section~\ref{sec:fitting}, we fit our model to the CFIRB angular power spectra and intensity. Section~\ref{sec:comparisons} shows comparisons between our model and other infrared observations. In Section~\ref{sec:implications}, we discuss the implications of our model, including the galaxy--halo connection and the cosmic star formation history; in Section~\ref{sec:discussion}, we discuss the limitations of our model and possible improvements. We summarize in Section~\ref{sec:summary}. Throughout this paper, we use a flat $\Lambda$CDM cosmology based on the {\em Planck} 2013 results \citep{Planck13cosmo}; $\Omega_M$ = 0.31; $\Omega_\Lambda$ = 0.69; $h$ = 0.67. We use the linear matter power spectrum at $z=0$ calculated by {\sc CAMB} \citep{Lewis00} with $\Omega_b h^2 = 0.022$; $\Omega_c h^2 = 0.12$; $n_s = 0.96 $; $A_s = 2.215\times10^{-9}$. When converting SFR to IR luminosity, we use $\LIR = {\rm SFR}/{K}$, where $K=1.7\times10^{-10}\ \Msun\ \rm yr^{-1}\ \Lsun^{-1}$ based on the Salpeter initial mass function \citep{Kennicutt98}.

\label{sec:discussion} In this work, we show that the gas regulator model provides a qualitative description for the CFIRB power spectra but is unable to produce all the details in observations. In this section, we discuss the limitations of our implementation of the gas regulator model and possible improvements. In our implementation, we assume that most of the parameters are time-independent and mass-independent, and we incorporate the mass-dependence in the mass-loading factor (equation \ref{eq:etaM}) and the extra time-dependence in SFR (equation \ref{eq:fz}). These parametrizations attempt to capture the effects of feedback, but they do not capture the detailed physics and thus cannot reproduce observations perfectly. The effective mass-loading factor $\eta$, the accretion of gas $\fga$ and stars ($1-\fga$), and the return fraction $R$ can all have non-trivial time and mass dependence. Capturing the time and mass dependence accurately would require hydrodynamic simulations or semi-analytic models. The limitations of the gas regulator model have also been demonstrated in the literature. For example, DM14 have shown that their fiducial model systematically under-estimates the specific SFR at $1 < z < 4$. In our work, we use a few free parameters but are still unable to fit the data perfectly. \cite{KrumholzDekel12} implemented a metal-dependent SFR to take into account the fact that at high redshift, low-mass galaxies tend to have low metallicities and are unable to sustain a cool gas reservoir. Therefore, the SFR for low-metallicity galaxies is suppressed and is lower than what we would expect from the gas accretion rate. In our model, this effect is mimicked by the high effective mass-loading factor for low-mass galaxies. Our model qualitatively captures such trend; however, in principle, the SFR should be modelled self-consistently given the metallicity and dust. Furthermore, in our model we assume that all galaxies follow a simple modified body SED with the dust temperature calculated by assuming thermal equilibrium. This assumption is too simplistic and may be the reason why we have significantly worse fit in the 217 GHz band. Our model does not include starburst galaxies, which contribute to $\sim$10 per cent of the cosmic SFR density at $z\sim 2$ \citep{Rodighiero11,Sargent12} and are expected to have negligible contribution to CFIRB \citep{Shang12,Bethermin13}. Including the starburst galaxies could increase the bright end of the luminosity functions and number counts. However, an extra component for starburst would boost the power spectrum in the same way as a higher gas accretion rate would, and breaking such degeneracy would require a joint fit to the bright end of the luminosity functions. We apply the gas regulator model of galaxy evolution to describe dusty star-forming galaxies across cosmic time. We fit the model to the CFIRB power spectra observed by {\em Planck}. We compare our model predictions with the total CFIRB intensity measured by {\em COBE}, the correlation between CFIRB and CMB lensing potential measured by {\em Planck}, the bolometric IR luminosity functions up to $z=4$ from {\em Herschel} and {\em Spitzer}, and the total number counts from {\em Herschel}. The implications of our model are summarized as follows: \begin{itemize} \item The CFIRB power spectra favour a strong clustering of FIR galaxies. At $z=0$ ($z=2$), the large-scale galaxy bias is equivalent to the bias of dark matter haloes of mass $10^{13}$ ($10^{12.5}$) $\Msun$. This galaxy bias is consistent with the correlation between CFIRB and CMB lensing potential. \item The luminosity--mass relation from our model indicates that for massive haloes, the IR luminosity is higher than expected from the SFR constrained by UV and optical. This result is consistent with the high galaxy bias we have found. This excess in IR luminosity for massive haloes may come from the dust heated by old stellar populations. \item In our model, the luminosity--mass relation for low-mass haloes is lower than expected from the SFR constrained by UV and optical. These low-mass galaxies tend to be inefficient in absorbing UV photons, and their FIR emissions can underestimate the true SFR. \item Our model under-predicts the bright source counts of {\em Herschel}, slightly under-predicts the differential CFIRB intensity of {\em Herschel} for $z<1$, and over-predicts the CFIRB power spectra of {\em Herschel} at small scales. \item The cosmic star formation history from our model agrees with the recent compilation of \citet{MadauDickinson14} at $z<2$ but shows an excess at higher redshift. In addition, the total dust mass density across cosmic time is consistent with the results from {\em Herschel} CFIRB at $z>1$, while it is lower than the results from IR luminosity functions at $z<1$. \item Compared with SMGs selected from ground-based surveys, the galaxies in our model tend to have higher dust temperature ($T_{\rm dust} \gtrsim 25$ at $z=0$ and increases with redshift) and lower dust mass. \end{itemize} Our theoretical framework provides a simple, physically-motivated way to compare different FIR observations. It can be generalized to compute the foreground for various intensity mapping experiments. Our framework will also be useful for optimizing the survey designs and strategies for future FIR surveys. For example, the next generation CMB experiments, such as PIXIE \citep{Kogut11} and CORE \citep{CORE15}, will provide larger frequency coverage and/or higher angular resolution and sensitivity than {\em Planck} and will be able to provide better measurements for the CFIRB anisotropies as well as individual sources. In \cite{Wu16c}, we investigate the constraining power from future CFIRB experiments. The Far-IR Surveyor, which is currently explored by NASA\footnote{http://asd.gsfc.nasa.gov/firs/}, will reveal many more properties of dusty star-forming galaxies. Table~\ref{tab:cor} shows the correlation matrix of these parameters. Fig.~\ref{fig:MCMC} shows the 1-D and 2-D posterior distributions from the MCMC chains, which use the {\sc corner} software \citep{corner}. Figs.~\ref{fig:CL_217_sensitivity} and \ref{fig:CL_857_sensitivity} show the sensitivity of the power spectra to the model parameters at 217GHz and 857 GHz. In each panel, we increase or decrease a parameter by $2\sigma$. We note that the parameter $\sigma$ only affects the shot noise. Since shot noise dominates at larger $k$ and at higher frequency, the impacts of $\sigma$ is the strongest at large $k$ at 857 GHz. \begin{table*} \centering \setlength{\tabcolsep}{0.5em} \begin{tabular}{c|cccccccccccccc} \hline \rule[-2mm]{0mm}{6mm}&$\eta_0$&$\alpha_1$&$\alpha_2$&$\beta$&$\sigma$&$\delta$\\\hline \rule[-2mm]{0mm}{6mm}$\eta_0$&1.00&-0.59&-0.88&-0.48&-0.07&0.99\\ \rule[-2mm]{0mm}{6mm}$\alpha_1$&-0.59&1.00&0.55&0.36&-0.47&-0.53\\ \rule[-2mm]{0mm}{6mm}$\alpha_2$&-0.88&0.55&1.00&0.50&0.26&-0.86\\ \rule[-2mm]{0mm}{6mm}$\beta$&-0.48&0.36&0.50&1.00&-0.05&-0.56\\ \rule[-2mm]{0mm}{6mm}$\sigma$&-0.07&-0.47&0.26&-0.05&1.00&-0.10\\ \rule[-2mm]{0mm}{6mm}$\delta$&0.99&-0.53&-0.86&-0.56&-0.10&1.00\\\hline \end{tabular} \caption{Correlation matrix for the model parameters.} \label{tab:cor} \end{table*} \begin{figure*} \centering \includegraphics[width=2\columnwidth]{plots/corner.png} \caption[MCMC]{The 68\% and 95\% constraints of our model parameters. The diagonal panels show the posterior distribution and the 68\% constraint of each parameter.} \label{fig:MCMC} \end{figure*} \begin{figure*} \centering \includegraphics[width=2\columnwidth]{plots/CL_217_sensitivity.pdf} \vspace{-0.7cm} \caption[]{Sensitivity of angular power spectra (217 GHz) to model parameters. In each panel, the solid and dash curves correspond to increasing and decreasing the model parameters by $2\sigma$.} \label{fig:CL_217_sensitivity} \end{figure*} \begin{figure*} \centering \vspace{-0.5cm} \includegraphics[width=2\columnwidth]{plots/CL_857_sensitivity.pdf} \caption[]{Same as Fig.~\ref{fig:CL_217_sensitivity} but for 857 GHz.} \label{fig:CL_857_sensitivity} \end{figure*}

1607.02546

1607

1607.05199_arXiv.txt

We present gravitational wave (GW) signal predictions from four 3D multi-group neutrino hydrodynamics simulations of core-collapse supernovae of progenitors with $11.2 M_\odot$, $20 M_\odot$, and $27 M_\odot$. GW emission in the pre-explosion phase strongly depends on whether the post-shock flow is dominated by the standing accretion shock instability (SASI) or convection and differs considerably from 2D models. SASI activity produces a strong signal component below $250 \, \mathrm{Hz}$ through asymmetric mass motions in the gain layer and a non-resonant coupling to the proto-neutron star (PNS). Both convection- and SASI-dominated models show GW emission above $250 \, \mathrm{Hz}$, but with considerably lower amplitudes than in 2D. This is due to a different excitation mechanism for high-frequency $l=2$ motions in the PNS surface, which are predominantly excited by PNS convection in 3D. Resonant excitation of high-frequency surface g-modes in 3D by mass motions in the gain layer is suppressed compared to 2D because of smaller downflow velocities and a lack of high-frequency variability in the downflows. In the exploding $20 M_\odot$ model, shock revival results in enhanced low-frequency emission due to a change of the preferred scale of the convective eddies in the PNS convection zone. Estimates of the expected excess power in two frequency bands suggests that second-generation detectors will only be able to detect very nearby events, but that third-generation detectors could distinguish SASI- and convection-dominated models at distances of $\mathord{\sim} 10 \, \mathrm{kpc}$.

Despite impressive progress during recent years, the explosion mechanism powering core-collapse supernovae is still not fully understood. For ordinary supernovae with explosion energies up to $\mathord{\sim}10^{51} \,\mathrm{erg}$, the prevailing theory is the delayed neutrino-driven mechanism (see \citealp{janka_12,burrows_13} for current reviews). In this scenario, the shock wave formed during the rebound (bounce) of the inner core initially stalls and only propagates out to a radius of $\mathord{\sim}150 \,\mathrm{km}$. The energy needed to revitalise the shock is provided by the partial re-absorption of neutrinos emitted from the proto-neutron star (PNS). Hydrodynamical instabilities operating behind the stalled shock front have been found to be crucial for the success of this scenario as they help to push the shock further out by generating large Reynolds stresses (or ``turbulent pressure'', see \citealp{burrows_95,murphy_12,couch_15,mueller_15a}) and transporting neutrino-heated material out from the gain radius, which then allows the material to be exposed to neutrino heating over a longer ``dwell time'' \citep{buras_06b,murphy_08b}. Moreover, if the instabilities lead to the formation of sufficiently large high-entropy bubbles, the buoyancy of these bubbles can become high enough to allow them to rise and expand continuously \citep{thompson_00,dolence_13,fernandez_15}. Two such instabilities have been identified in simulations, namely the more familiar phenomenon of convection driven by the unstable entropy gradient arising due to neutrino heating \citep{bethe_90,herant_94,burrows_95,janka_96,mueller_97}, and the so-called standing accretion shock instability (SASI), which manifests itself in large-scale sloshing and spiral motions of the shock \citep{blondin_03,blondin_06,foglizzo_07,ohnishi_06,ohnishi_08,scheck_08,guilet_12,foglizzo_15}. After initial setbacks in three-dimensional (3D) supernova modelling, we are now starting to see the emergence of the first generation of successful 3D simulations of explosions with three-flavour multi-group neutrino transport, culminating in the recent models of the Garching and Oak~Ridge groups \citep{melson_15a,melson_15b,lentz_15} with their rigorous treatment of the transport and neutrino microphysics in addition to many more obtained with more approximate transport schemes, as for example the studies of \citet{takiwaki_12,takiwaki_14}, \citet{mueller_15b} and \citet{roberts_16}.\footnote{ \citet{takiwaki_12,takiwaki_14} employ the isotropic diffusion source approximation \citep{liebendoerfer_09} and use further approximations to treat heavy lepton neutrinos. \cite{takiwaki_14} employ a leakage scheme to account for heavy lepton neutrinos and \cite{takiwaki_12} neglect the effect of these neutrinos altogether. \cite{mueller_15b} utilises the stationary fast multi-group transport scheme of \cite{mueller_15a}, which at high optical depths solves the Boltzmann equation in a two-stream approximation and matches the solution to an analytic variable Eddington factor closure at low optical depths. \citet{roberts_16} employ a full 3D two-moment (M1) solver in general relativistic simulations, but ignore velocity-dependent terms.} Our means to validate these numerical models by observations are limited. Photon-based observations of supernovae and their remnants (e.g.\ mixing in the envelope, see \citealp{wongwathanarat_15} and references therein; pulsar kicks, \citealp{scheck_06,wongwathanarat_10b,wongwathanarat_13,nordhaus_12}) provide only indirect constraints on the workings of the hydrodynamic instabilities in the inner engine of a core-collapse supernova. For a nearby, Galactic supernova event, messengers from the core in the form of neutrinos and gravitational waves (GWs) could furnish us with a direct glimpse at the engine. Neutrinos, for example, could provide a smoking gun for SASI activity through fast temporal variations \citep{marek_08,lund_10,brandt_11,tamborra_13,tamborra_14b,mueller_14} and could even allow a time-dependent reconstruction of the shock trajectory \citep{mueller_14}. Likewise, a detection of GWs could potentially help to unveil the multi-dimensional effects operating in the core of a supernova. The signal from the collapse and bounce of rapidly rotating iron cores and triaxial instabilities in the early post-bounce phase has long been studied in 2D (i.e.\ under the assumption of axisymmetry) and 3D (e.g.\ \citealp{ott_06_a,dimmelmeier_07_a,dimmelmeier_08,scheidegger_08,abdikamalov_10}). Understanding the GW signal generated by convection and the SASI in the more generic case of slow or negligible rotation has proved more difficult due to a more stochastic nature of the signal. During the recent years, however, a coherent picture of GW emission has emerged from parameterised models \citep{murphy_09} and first-principle simulations of supernova explosions in 2D \citep{marek_08,mueller_13}: The models typically show an early, low-frequency signal with typical frequencies of $\mathord{\sim} 100 \,\mathrm{Hz}$ arising from shock oscillations that are seeded by prompt convection \citep{marek_08,murphy_09,yakunin_10,mueller_13,yakunin_15}. This signal component is followed by a high-frequency signal with stochastic amplitude modulations that is generated by forced oscillatory motions in the convectively stable neutron star surface layer \citep{marek_08,murphy_09,mueller_13} with typical frequencies of $300 \ldots 1000 \,\mathrm{Hz}$ that closely trace the Brunt-V\"ais\"ala frequency in this region \citep{mueller_13}. Prior to the explosion, these oscillations, tentatively identified as $l=2$ surface g-modes by \citet{mueller_13}, are primarily driven by the downflows impinging onto the neutron star, whereas PNS convection takes over as the forcing agent a few hundred milliseconds after shock revival as accretion dies down. This high-frequency contribution dominates the energy spectrum and the total energy emitted in GWs can reach $\mathord{\sim} 10^{46} \,\mathrm{erg}$ \citep{mueller_13,yakunin_15}. Since 3D supernova models have proved fundamentally different to 2D models in many respects, it stands to reason that much of what we have learned about GW emission from first-principle 2D models will need to be revised. In 2D, the inverse turbulent cascade \citep{kraichnan_76} facilitates the emergence of large-scale flow structures also in convectively-dominated models and helps to increase the kinetic energy in turbulent fluid motions in the post-shock region \citep{hanke_12}. Furthermore, accretion downflows impact the PNS with much higher velocities in 2D than in 3D \citep{melson_15a} due to the inverse turbulent cascade and the stronger inhibition of Kelvin-Helmholtz instabilities at the interface of supersonic accretion downflows \citep{mueller_15b}. In the SASI-dominated regime, on the other hand, the additional dimension allows the development of the spiral mode \citep{blondin_07a,blondin_07b,fernandez_10} in 3D, which can store more non-radial kinetic energy than pure sloshing motions in 2D \citep{hanke_13,fernandez_15}, contrary to earlier findings of \cite{iwakami_08}. Such far-reaching differences between 2D and 3D cannot fail to have a significant impact on the GW signal. While the impact of 3D effects on the GW signals from the post-bounce phase has been investigated before, all available studies have relied on a rather approximate treatment of neutrino heating and cooling such as simple light-bulb models \citep{mueller_97,kotake_09,kotake_11}, grey neutrino transport \citep{fryer_04,mueller_e_12}, or a partial implementation of the isotropic diffusion source approximation of \citet{liebendoerfer_09} in the works of \citet{scheidegger_08,scheidegger_10}, which were also limited to the early post-bounce phase. Arguably, none of these previous studies have as yet probed precisely the regimes encountered by the best current 3D simulations (e.g.\ the emergence of a strong SASI spiral mode) and therefore cannot be relied upon for quantitative predictions of GW amplitudes and spectra, which are extremely sensitive to the nature of hydrodynamic instabilities, the neutrino heating, and the contraction of the PNS. In this paper, we present GW waveforms of the first few hundred milliseconds of the post-bounce phase computed from 3D models with multi-group neutrino transport. Waveforms have been analysed for four supernova models of progenitors with zero-age main sequence (ZAMS) masses of $11.2 M_\odot$, $20 M_\odot$ (for which we study an exploding and a non-exploding simulation), and $27 M_\odot$. With four simulations based on these three different progenitors, we cover both the convective regime ($11.2 M_\odot$) and the SASI-dominated regime ($20 M_\odot$, $27 M_\odot$). Our aim in studying waveforms from these progenitors is twofold: On the one hand, we shall attempt to unearth the underlying hydrodynamical phenomena responsible for the GW emission in different regions of the frequency spectrum during different phases of the evolution. We shall also compare the GW emission in 3D and 2D models, which will further illuminate dynamical differences between 2D and 3D. Furthermore, with 3D models now at hand, we are in a position to better assess the detectability of GWs from the post-bounce phase in present and future instruments than with 2D models affected by the artificial constraint of axisymmetry. One of our key findings is that the GW signal from SASI-dominated models is clearly differentiated from convection-dominated model by strong emission in a low-frequency band around $100 \ldots 200 \, \mathrm{Hz}$. Very recently, \cite{kuroda_16} also studied the GW signal features (in models using grey neutrino transport) during phases of SASI activity for a $15 M_\odot$ star, comparing results for three different nuclear equations of state. Going beyond \cite{kuroda_16}, we clarify why this signature has not been seen in 2D models and point out that the hydrodynamic processes underlying this low-frequency signal are quite complex and seem to require a coupling of SASI motions to deeper layers inside the PNS. Moreover, we show that broadband low-frequency GW emission can also occur after the onset of the explosion and is therefore not an unambiguous signature of the SASI. We also provide a more critical assessment of the detectability of this new signal component, suggesting that it may only be detectable with second-generation instruments like Advanced LIGO for a very nearby event at a distance of $2 \, \mathrm{kpc}$ or less. Our paper is structured as follows: We first give a brief description of the numerical setup and the extraction of GWs in Section~\ref{sec:numerics}. In Section~\ref{sec:structure}, we present a short overview of the GW waveforms and then analyse the hydrodynamical processes contributing to different parts of the spectrum in detail. We also compare our results to recent studies based on 2D first-principle models. In Section~\ref{sec:obs}, we discuss the detectability of the predicted GW signal from our three progenitors by Advanced LIGO \citep{advligo_15}, and by the Einstein Telescope \citep{et_12} as next-generation instrument. We also comment on possible inferences from a prospective GW detection. Our conclusions and a summary of open questions for future research are presented in Section~\ref{sec:con}.

\label{sec:con} We have studied the GW signal from the accretion phase and the early explosion phase of core-collapse supernovae based on four recent 3D multi-group neutrino hydrodynamics simulations. We considered four models based on three progenitors with ZAMS masses of $11.2 M_\odot$, $20 M_\odot$, and $27 M_\odot$. The three non-exploding models enabled us to study the phase between bounce and shock revival. We covered both the SASI-dominated regime (model s20, \citealp{tamborra_14b}; model s27, \citealp{hanke_13}), as well as the convection-dominated regime (model s11.2, \citealp{tamborra_14a}). Additionally, the exploding $20 M_\odot$ model s20s \citep[][with a modified axial-vector coupling constant for neutral current scattering]{melson_15b} illustrates changes in the GW signal in exploding models. Since our treatment of the microphysics and the neutrino transport is on par with previous works on the GW signal from 2D simulations \citep{marek_08,yakunin_10,mueller_13,yakunin_15}, we were in the position to conduct a meaningful comparison of GW emission in 2D and 3D during the accretion and explosion phase for the first time. To this end, we included the $27 M_\odot$ 2D models of \citet{mueller_12b} and \citet{hanke_13} in our study. Our analysis showed differences between the GW emission in 2D and 3D. The prominent, relatively narrow-banded emission at high-frequencies that is characteristic of 2D models is significantly reduced. With the reduction of the high-frequency emission, distinctive broadband \emph{low-frequency} emission in the range between $100 \, \mathrm{Hz}$ and $200 \, \mathrm{Hz}$ emerges as a characteristic feature during episodes of SASI activity and during the explosion phase of model s20s. The low-frequency emission does also exist in the 2D models, but it is completely overwhelmed by the high-frequency emission. This conclusion is somewhat model dependent, because in one of our 2D models, s27-2D, high-frequency GW emission is low and the low-frequency component becomes very prominent. We discussed these differences extensively from two vantage points: On the one hand, we investigated the underlying hydrodynamic processes responsible for GW emission and showed how the changes in the GW signal in 3D are related to critical differences in flow dynamics in 3D compared to 2D. On the other hand, we outlined the repercussions of these changes for future GW observations and sketched possible inferences that could be drawn from the detection of a Galactic event by third-generation instruments. With regard to the hydrodynamic processes responsible for GW emission, our findings can be summarised as follows: \begin{enumerate} \item There is a high-frequency signal component that closely traces the buoyancy frequency in the PNS surface region in 2D and 3D, i.e.\ the roughly isothermal atmosphere layer between the PNS convection zone and the gain region acts as frequency stabiliser for forced oscillatory motions in both cases. However, the high-frequency component mostly stems from aspherical mass motions in and close to the overshooting region of PNS convection in 3D, whereas it stems from mass motion close to the gain radius in 2D. This indicates that quasi-oscillatory mass motions at high frequencies are instigated \emph{only by PNS convection in 3D} even during the pre-explosion phase, whereas forcing by the SASI and convection in the gain region is dominant in 2D. The resulting \emph{amplitudes of the high-frequency component are considerably lower in 3D than in 2D}. \item We ascribe the strong excitation of high-frequency surface g-mode oscillations in 2D to several causes: The inverse turbulent cascade in 2D leads to larger impact velocities of the downflows and creates large flow structures that can effectively excite $l=2$ oscillations that give rise to GW emission. Braking of downflows by the forward turbulent cascade and fragmentation into smaller eddies strongly suppress surface g-mode excitation in 3D. Moreover, the spectrum of turbulent motions does not extend to high frequencies in 3D both in SASI-dominated and convection-dominated models so that the resonant excitation of the $l=2$ surface g-mode at its eigenfrequency becomes ineffective. \item In 3D, low-frequency GW emission in the pre-explosion phase ultimately stems from the global modulation of the accretion flow by the SASI. Because of frequency doubling and/or the contribution from the $l=2$ mode, the typical frequencies of this component are of the order of $100 \ldots 200 \, \mathrm{Hz}$, i.e.\ somewhat higher than the typical frequency of the $l=1$ modes of the SASI. Mass motions in the post-shock region, the PNS surface region and the PNS convection zone all contribute to this low-frequency component, which indicates that the modulation of the accretion flow is still felt deep below the gain radius as the accreted matter settles down onto the PNS. Moreover, our analysis of the detection prospects shows that \emph{the low-frequency component of the signal at $\mathord{\gtrsim} 100 \, \mathrm{Hz}$ becomes a primary target in terms of detectability} in contrast to previous 2D results. \item By contrast, convective models characterised by mass motions of intermediate- and small-scale like s11.2 show very little GW emission at low frequencies. The high-frequency emission, on the other hand, is excited primarily by PNS convection and is therefore less sensitive to the dominant instability (convection or SASI) in the post-shock region. \emph{Thus, the ratio of high-frequency to low-frequency GW power can potentially be used to distinguish SASI- and convection-dominated models in the pre-explosion phase.} \item However, strongly enhanced low-frequency emission can also occur due to a change of the preferred scale of the convective eddies in the PNS convection zone as exemplified by model s20s, where the dominant mode shifts from $l=1$ to $l=2$ late in the simulation. Since this does not occur in the corresponding non-exploding model s20, one can speculate that this behaviour is due to changes in the accretion flow and neutron star cooling associated with shock revival. If this behaviour is generic for exploding models enhanced GW emission may still remain a fingerprint of shock revival as it is in 2D \citep{murphy_09,mueller_13}. With only one explosion model available to us, this conclusion does not rest on safe ground; more 3D explosion models are needed to check whether enhanced low-frequency GW emission after shock revival is indeed a generic phenomenon. \end{enumerate} It is obviously of interest whether future GW observations will be able to discriminate between models with such distinctively different behaviour as the ones presented here. Without an elaborate statistical analysis, only limited conclusions can be drawn concerning this point. In this paper, we confined ourselves to rough estimates based on the expected excess power in second- and third-generation GW detectors in two bands at low ($20 \ldots 250 \, \mathrm{Hz}$) and high ($250 \ldots 1200 \, \mathrm{Hz}$) frequency. Due to the reduction of the signal amplitudes compared to 3D, the prospects for second-generation detectors appear rather bleak; even the SASI-dominated models s20, s20s, and s27 could not be detected out further than $\mathord{\sim} 2 \, \mathrm{kpc}$ with AdvLIGO at a confidence level of $95\%$. Third-generation instruments like the Einstein Telescope, however, could not only detect all of our models at the typical distance of a Galactic supernova ($\mathord{\sim} 10 \, \mathrm{kpc}$) and strong GW emitters like s20s out to $50 \, \mathrm{kpc}$; the expected signal-to-noise ratios could even be high enough to distinguish models with enhanced low-frequency emission due to SASI from convective models based on the ``colour'' of the GW spectrum. In conjunction with timing information and the neutrino signal, it may also be possible to distinguish enhanced low-frequency emission from the SASI from enhanced GW emission after shock revival as in model s20s. However, more work is obviously needed to fully exploit the potential of GWs as a probe of the supernova engine in the case of ``ordinary'', slowly rotating supernovae for which PNS convection and the SASI are the dominant sources of GW emission. Desiderata for the future include a much broader range of 3D explosion models to determine to what extent the aforementioned features in the GW signal are generic. With waveforms from longer explosion simulations, the prospects for detecting a Galactic supernova in GWs with second generation instruments may also appear less bleak than they do now based on our biased selection that includes only one explosion model evolved to $200 \, \mathrm{ms}$ after shock revival. Furthermore, it is conceivable that much more information can be harvested from the GW signals than our simple analysis suggests. Several authors \citep{logue_12,hayama_15,gossan_15} have already demonstrated the usefulness of a powerful statistical machinery in assessing the detectability of supernovae in GWs and distinguishing different waveforms (e.g.\ from rotational collapse and hot-bubble convection, \citealp{logue_12}). Peeling out the more subtle differences between SASI- and convection dominated models from GW signals in the face of greatly reduced signal amplitudes certainly presents a greater challenge, but third-generation instruments will nonetheless make it an effort worth undertaking. The GW analysis presented in this work is based on three non-rotating progenitors, {\comment and it remains to be seen whether the findings from these simulations are generic. For GW detection, it is particularly important to ascertain whether the overall reduction of the signal from SASI and convection in 3D compared to 2D is always as strong as in our models, where the difference is a factor of $\mathord{\sim} 10$. This has recently been questioned by \citet{yakunin_17}, who reported considerably higher amplitudes for a $15 M_\odot$ progenitor than in our models and found the energy emitted in GW to be similar in their 2D and 3D simulations. Considering that we obtain weaker GW signals in 3D in models that probe a variety of different regimes, and that other 3D studies \citep{mueller_e_12,kuroda_16} predict amplitudes in line with our findings (albeit with less rigorous neutrino transport and without a 2D/3D comparison) suggests that small amplitudes $|A|\lesssim 5 \, \mathrm{cm}$ are generic in 3D and that the strong amplitudes in \citet{yakunin_17} are the exception rather than the norm and need further explanation. Nonetheless, the range of variation in GW amplitude in 3D deserves to be explored further in the future. There are various properties of the pre-collapse cores that will (or at least could) impact the GW signal. The influence of rotation is well known:} In rapidly rotating models there is a strong GW burst associated with the rebound of the core \citep{mueller_82}. During the post-bounce phase rotation can lead to a bar-like deformation of the core \citep{rampp_98,shibata_05} or the development of low-mode spiral instabilities \citep{ott_05,kuroda_14,takiwaki_16}. These flow patterns in turn lead to strong GW emission at frequencies determined by the rotational frequency. In addition, rotation can modulate processes already present in nonrotating models, for example prompt convection or the SASI. In the models presented by \citet{dimmelmeier_08} and \citet{ott_12} only models with moderate rotation rates (and nonrotating models) exhibit prompt convection. The coupling between rotation and SASI activity can lead to an enhanced growth rate of the spiral SASI mode \citep{blondin_07a,yamasaki_08,iwakami_09,kazeroni_16,janka_16}. Whether a significant proportion of supernova progenitors have moderately rotating (let alone rapidly rotating) cores is unclear. Stellar evolution models that include the effects of magnetic fields predict rather slowly rotating pre-collapse cores \citep{heger_05}. Furthermore, the angular momentum loss due to stellar winds seems to be underestimated by stellar evolution models, compared to results from asteroseismology \citep{cantiello_14}. Predictions of the initial rotation rate of pulsars, based on their current spin-down rate and age, suggest that a large fraction of the pulsar population is born with rotation periods of the order of tens to hundreds of milliseconds \citep{popov_12,noutsos_13}. There is also the issue of starting the simulations from spherically symmetric progenitor models. {\comment Current GW predictions like ours rely on explicitly imposed (this study) or numerical seed perturbations to trigger the development of non-radial instabilities, and it needs to be explored further whether the level of seed perturbations is partly responsible for differences in the GW amplitudes calculated by different groups (e.g.\ this study and \citealt{yakunin_17}). Moreover, it} has been found that {\comment physical seed} asymmetries in the burning shells of the progenitor can influence the shock dynamics and even help to ensure a successful explosion \citep{burrows_96,fryer_04,arnett_11,couch_13,mueller_15a}. Any change in the initial conditions that leads to a significant change in the dynamics of the supernova core should be expected to impact the GW signal. Therefore, it will be important to keep improving the predicted GW signals, in hand with the improvement of core collapse models.

1607.05199

1607

1607.02400_arXiv.txt

In a new classification of merging binary neutron stars (NSs) we separate short gamma-ray bursts (GRBs) in two sub-classes. The ones with $E_{\rm iso}\lesssim10^{52}$~erg coalesce to form a massive NS and are indicated as short gamma-ray flashes (S-GRFs). The hardest, with $E_{\rm iso}\gtrsim10^{52}$~erg, coalesce to form a black hole (BH) and are indicated as genuine short-GRBs (S-GRBs). Within the fireshell model, S-GRBs exhibit three different components: the P-GRB emission, observed at the transparency of a self-accelerating baryon-$e^+e^-$ plasma; the prompt emission, originating from the interaction of the accelerated baryons with the circumburst medium; the high-energy (GeV) emission, observed after the P-GRB and indicating the formation of a BH. GRB 090510 gives the first evidence for the formation of a Kerr BH or, possibly, a Kerr-Newman BH. Its P-GRB spectrum can be fitted by a convolution of thermal spectra whose origin can be traced back to an axially symmetric dyadotorus. A large value of the angular momentum of the newborn BH is consistent with the large energetics of this S-GRB, which reach in the $1$--$10000$~keV range $E_{\rm iso}=(3.95\pm0.21)\times10^{52}$~erg and in the $0.1$--$100$~GeV range $E_{\rm LAT}=(5.78\pm0.60)\times10^{52}$~erg, the most energetic GeV emission ever observed in S-GRBs. The theoretical redshift $z_{\rm th}=0.75\pm0.17$ that we derive from the fireshell theory is consistent with the spectroscopic measurement $z=0.903\pm0.003$, showing the self-consistency of the theoretical approach. All S-GRBs exhibit GeV emission, when inside the \textit{Fermi}-LAT field of view, unlike S-GRFs, which never evidence it. The GeV emission appears to be the discriminant for the formation of a BH in GRBs, confirmed by their observed overall energetics.

Thanks to a fortunate coincidence of observations by AGILE, \textit{Fermi}, and \textit{Swift} satellites, together with the optical observations by the VLT/FORS2 and the Nordic Optical Telescope, it has been possible to obtain an unprecedented set of data, extending from the optical-UV, through the X-rays, all the way up to the high energy (GeV) emission, which allowed detailed temporal/spectral analyses on GRB 090510 \citep{Depasquale2010}. In contrast with this outstanding campaign of observations, a theoretical analysis of the broadband emission of GRB 090510 has been advanced within the synchrotron/self-synchrotron Compton (SSC) and traditional afterglow models \citep[see, e.g., sections 5.2.1 and 5.2.2 in][]{Ackermann2010}. Paradoxically, this same methodology has been applied in the description of markedly different type of sources: e.g., \citet{Soderberg} for the low energetic long GRB 060218, \citet{2014ApJ...781...37P} for the high energetic long GRB 130427A, and \citet{2006ApJ...650..261S} for the S-GRF 051221A \citep[see also][and references therein]{2008A&A...487..533C}. In the meantime, it has become evident that GRBs can be subdivided into a variety of classes and sub-classes \citep{Wang2015,Muccino2015,Ruffini2016}, each of them characterized by specific different progenitors which deserve specific theoretical treatments and understanding. In addition every sub-class shows different episodes corresponding to specifically different astrophysical processes, which can be identified thanks to specific theoretical treatments and data analysis. In this article, we take GRB 090510 as a prototype for S-GRBs and perform a new time-resoved spectral analysis, in excellent agreement with the above temporal and spectral analysis performed by, e.g., the \textit{Fermi} team. Now this analysis, guided by a theoretical approach successfully tested in this new family of S-GRBs \citep{Muccino2013,Muccino2015}, is directed to identify a precise sequence of different events made possible by the exceptional quality of the data of GRB 090510. This include a new structure in the thermal emission of the P-GRB emission, followed by the onset of the GeV emission linked to the BH formation, allowing, as well, to derive the structure of the circumburst medium from the spiky structure of the prompt emission. This sequence, for the first time, illustrates the formation process of a BH. Already in February 1974, soon after the public announcement of the GRB discovery \citep{Strong1975}, \citet{DamourRuffini1975} presented the possible relation of GRBs with the vacuum polarization process around a Kerr-Newman BH. There, evidence was given for: a) the formation of a vast amount $e^+e^-$-baryon plasma; b) the energetics of GRBs to be of the order of $E_{\rm max}\approx10^{54} M_{\rm BH}/M_\odot$~erg, where $M_{\rm BH}$ is the BH mass; c) additional ultra-high energy cosmic rays with energy up to $\sim10^{20}$~eV originating from such extreme process. A few years later, the role of an $e^+e^-$ plasma of comparable energetics for the origin of GRBs was considered by \citet{CavalloRees} and it took almost thirty years to clarify some of the analogies and differences between these two processes leading, respectively, to the alternative concepts of ``fireball" and ``fireshell" \citep{Aksenov2007,Aksenov2009}. In this article we give the first evidence for the formation of a Kerr Newman BH, in GRB 090510, from the merger of two massive NSs in a binary system. GRBs are usually separated in two categories, based on their duration properties \citep[e.g.][]{1981Ap&SS..80....3M,Dezalay1992,Klebesadel1992,Kouveliotou1993,Tavani1998}. Short GRBs have a duration $T_{90} \lesssim 2$ s while the remaining ones with $T_{90} \gtrsim 2$ s are traditionally classified as long GRBs. Short GRBs are often associated to NS-NS mergers (see e.g.~\citealt{Goodman1986,Paczynski1986,Eichler1989, Narayan1991, Meszaros1997,Rosswog2003,Lee2004,2007PhR...442..166N,2016arXiv160403445E,2016ApJ...824L...6R}; see also \citealt{Berger2014} for a recent review): their host galaxies are of both early- and late-type, their localization with respect to the host galaxy often indicates a large offset \citep{Sahu1997,vanParadijs1997,Bloom2006,Troja2008,Fong2010,Berger2011,Kopac2012} or a location of minimal star-forming activity with typical circumburst medium (CBM) densities of $\sim10^{-5}$--$10^{-4}$ cm$^{-3}$, and no supernovae (SNe) have ever been associated to them. The progenitors of long GRBs, on the other hand, have been related to massive stars \citep{WoosleyBloom2006}. However, in spite of the fact that most massive stars are found in binary systems \citep{Smith2014}, that most type Ib/c SNe occur in binary systems \citep{Smith2011} and that SNe associated to long GRBs are indeed of type Ib/c \citep{DellaValle2011}, the effects of binarity on long GRBs have been for a long time largely ignored in the literature. Indeed, until recently, long GRBs have been interpreted as single events in the jetted \textit{collapsar} fireball model (see e.g.~\citealt{Woosley1993,ReesMeszaros1992,Kobayashi1997,Piran2005,Gehrels2009,KumarZhang2015} and references therein). Multiple components evidencing the presence of a precise sequence of different astrophysical processes have been found in several long GRBs (e.g.~\citealt{Izzo2012}, \citealt{Penacchioni2012}). Following this discovery, further results led to the introduction of a new paradigm expliciting the role of binary sources as progenitors of the long GRB-SN connection. New developments have led to the formulation of the Induced Gravitational Collapse (IGC) paradigm \citep{Ruffini2001a,Ruffini2007,Rueda2012,Wang2015}. The IGC paradigm explains the GRB-SN connection in terms of the interactions between an evolved carbon-oxygen core (CO$_{\rm core}$) undergoing a SN explosion and its hypercritical accretion on a binary NS companion \citep{Ruffini2015}. The large majority of long bursts is related to SNe and are spatially correlated with bright star-forming regions in their host galaxies \citep{Fruchter2006,Svensson2010} with a typical CBM density of $\sim1$ cm$^{-3}$ \citep{Izzo2012,Penacchioni2012}. A new situation has occurred with the observation of the high energy GeV emission by the \textit{Fermi}-LAT instrument and its correlation with both long and short bursts with isotropic energy $E_{\mathrm{iso}} \gtrsim 10^{52}$ erg, which has been evidenced in \cite{Wang2015} and \cite{Muccino2015}, respectively. On the basis of this correlation the occurrence of such prolonged GeV emission has been identified with the onset of the formation of a BH \citep{Wang2015,Muccino2015}. As recalled above, the long GRBs associated to SNe have been linked to the hypercritical accretion process occurring in a tight binary system when the ejecta of an exploding CO$_{\rm core}$ accretes onto a NS binary companion \citep[see, e.g.,][]{Rueda2012,Fryer2014,Becerra}. When the hypercritical accretion occurs in a widely separated system with an orbital separation $>10^{11}$~cm \citep{Becerra}, the accretion is not sufficient to form a BH. For these softer systems with rest-frame spectral peak energy $E_{\rm peak}<200$ keV the upper limit of their observed energy is $E_{\mathrm{iso}}\approx10^{52}$ erg, which corresponds to the maximum energy attainable in the accretion onto a NS \citep{Wang2015}. Such long a burst corresponds to an X-ray flash (XRF). The associated X-ray afterglow is also explainable in terms of the interaction of the prompt emission with the SN ejecta (Fryer et al., in preparation). In these systems no GeV emission is expected in our theory and, indeed, is not observed. Interestingly, a pioneering evidence for such an X-ray flash had already been given in a different context by \cite{Heise2003}, \cite{Amati2004}, and \cite{Soderberg}. For tighter binaries ($<10^{11}$~cm, \citealt{Becerra}), the hypercritical accretion onto the companion NS leads to the formation of a BH. For these harder systems with $E_{\rm peak}>200$ keV the lower limit of their observed energy is $E_{\mathrm{iso}}\approx10^{52}$ erg, which necessarily needs the accretion process into a BH. An associated prolonged GeV emission occurs after the P-GRB emission and at the beginning of the prompt emission, and originates at the onset of the BH formation \citep{Wang2015}. These more energetic events are referred to as binary-driven hypernovae (BdHNe). Specific constant power-law behaviors are observed in their high energy GeV, X-rays, and optical luminosity light curves \citep{Pisani2013,Ruffini2014,Wang2015}. In total analogy, the formation of a BH can occur in short bursts, depending on the mass of the merged core of the binary system. When the two NS masses are large enough, the merged core can exceed the NS critical mass and the BH formation is possible. In the opposite case, a massive NS (MNS) is created, possibly, with some additional orbiting material to guarantee the angular momentum conservation. We then naturally expect the existence of two short bursts sub-classes: authentic short GRBs (S-GRBs), characterized by the formation of a BH \citep{Muccino2015}, with $E_{\mathrm{iso}} \gtrsim 10^{52}$ erg, a harder spectrum (see section \ref{epeakeiso}) and associated with a prolonged GeV emission (see section \ref{GeVemission}); short gamma-ray flashes (S-GRFs), producing a MNS \citep{Muccino2015}, with $E_{\mathrm{iso}} \lesssim 10^{52}$ erg. In this second sub-class, of course, the GeV emission should not occur and, indeed, is never observed. \begin{figure} \centering \includegraphics[width=\hsize,clip]{f1a} \caption{Projection of the dyadotorus of a Kerr-Newman BH corresponding to selected values of the ratio $E/E_c$, where $E_c$ is the critical value for vacuum polarization and $E$ is the electric field strength. The plot assumes a black hole mass energy $\mu = M_{\mathrm{BH}} / M_\odot = 10$. Figure reproduced from \cite{Cherubini2009} with their kind permission.} \label{dyadotorus} \end{figure} Following the discovery of the first prototype of this S-GRB class, namely GRB 090227B \citep{Muccino2013}, the first detailed analysis of such a genuine short GRB originating from a binary NS merger leading to a BH was done for GRB 140619B by \cite{Muccino2015}, determining as well the estimated emission of gravitational waves. The latter has been estimated following the method applied by \cite{Oliveira2014} for GRB 090227B. From the spectral analysis of the early $\sim 0.2$ s, they inferred an observed temperature $kT = (324 \pm 33)$ keV of the $e^+e^-$ plasma at transparency (P-GRB), a theoretically derived redshift $z = 2.67 \pm 0.37$, a total burst energy $E^{\mathrm{tot}}_{e^+e^-} = (6.03 \pm 0.79) \times 10^{52}$ erg, a rest-frame peak energy $E_{p,i} = 4.7$ MeV, a baryon load $B = (5.52 \pm 0.73) \times 10^{-5}$, and an average CBM density $n_{\mathrm{CBM}} = (4.7 \pm 1.2) \times 10^{-5}$ cm$^{-3}$. We turn in this article to the most interesting case of GRB 090510 which has, in addition to very similar properties of the members of this new class of S-GRB sources, a spectroscopically determined value of the redshift and represents one of the most energetic sources of this family both in the $\gamma$-ray and in the GeV ranges. Actually, a first attempt to analyze GRB 090510 was made by interpreting this source as a long GRB \citep{Muccino2013090510}. An unusually large value of the CBM density was needed in order to fit the data: this interpretation was soon abandoned when it was noticed that GRB 090510 did not fulfill the nesting conditions of the late X-ray emission typical of long GRBs \citep{Ruffini2014}, see also section \ref{xray} and Figure \ref{episode3}. In light of the recent progress in the understanding of the fireshell theory, we address the interpretation of GRB 090510 as the merging of a binary NS. We give clear evidence for the validity of this interpretation. In view of the good quality of the data both in $\gamma$- rays and in the GeV range, we have performed a more accurate description of the P-GRB, best fitted by a convolution of thermal spectra. This novel feature gives the first indication for the existence of an axially symmetric configuration of the dyadotorus emitting the $e^+e^-$ plasma which had been previously theoretically considered and attentively searched for. This gives the first indication that indeed the angular momentum plays a role and a dyadotorus is formed, as theoretically predicted in a series of papers \citep[see][and Figure \ref{dyadotorus}]{Cherubini2009,Ruffini2009}. This naturally leads to the evidence for the formation of a rotating BH as the outcome of the gravitational collapse. We turn then to the main new feature of GRB 090510 which is the high energy $0.1$--$100$ GeV emission (see Figure \ref{090510GeV}). The direct comparison of the GeV emission in this source and in the BdHNe 130427A shows the remarkable similarities of these two GeV components (see Figure \ref{090510GeV}). The fact that the S-GRB 090510 originates from a binary NS merger and the BdHN 130427A from the IGC of a SN hypercritical accretion process onto a companion NS clearly points to the BH as originating this GeV emission, the reason being that these two astrophysical systems are different in their progenitors and physical process and have in the formation of a BH their unique commonality. This paper is structured as follows: in section 2 we summarize the relevant aspects of the fireshell theory and compare and contrast it with alternative approaches. In section 3 we discuss the recent progress on the NS equilibrium configuration relevant for S-GRBs and BdHNe. In section 4 we move on to describe the observations of GRB 090510 and their analysis. The S-GRB nature of GRB 090510 is justified in section 5, and we offer an interpretation of our results in section 6. Section 7 concludes this work. A standard flat ${\Lambda}$CDM cosmological model with ${\Omega}_m = 0.27$ and $H_0 = 71$ km s$^{-1}$ Mpc$^{-1}$ is adopted throughout the paper.

It is interesting to recall some of the main novelties introduced in this paper with respect to previous works on GRB 090510. Particularly noteworthy are the differences from the previous review of short bursts by \citet{2007PhR...442..166N}, made possible by the discovery of the high energy emission by the Fermi team in this specific source \citep{Ackermann2010}. A new family of short bursts characterized by the presence of a BH and associated high energy emission when LAT data are now available, comprehends GRBs 081024B, 090227B, 090510, 140402A, and 140619B (see, e.g., Figure \ref{090510GeV}). The excellent data obtained by the Fermi team and interpreted within the fireshell model has allowed to relate in this paper the starting point of the high energy emission with the birth of a BH. Our fireshell analysis assumes that the $\gamma$-ray and the GeV components originate from different physical processes. First, the interpretation of the prompt emission differs from the standard synchrotron model: we model the collisions of the baryon accelerated by the GRB outflow with the ambient medium following a fully relativistic approach (see Section 2). Second, we assume that the GeV emission originates from the matter accretion onto the newly-born BH and we show that indeed the energy requirement is fulfilled. This approach explains also the delayed onset of the GeV emission, i.e., it is observable only after the transparency condition, namely after the P-GRB emission. The joint utilization of the excellent data from the \textit{Fermi}-GBM NaI-n6 and n7 and the BGO-b1 detectors and from the Mini-Calorimeter on board AGILE \citep{Giuliani2010} has given strong observational support to our theoretical work. GRB 090510 has been analyzed in light of the recent progress achieved in the fireshell theory and the resulting new classification of GRBs. We show that GRB 090510 is a S-GRB, originating in a binary NS merger (see figure~\ref{rt_gw}). Such systems, by the absence of the associated SN events, are by far the simplest GRBs to be analyzed. Our analysis indicates the presence of three distinct episodes in S-GRBs: the P-GRB, the prompt emission, and the GeV emission. By following the precise identification of successive events predicted by the fireshell theory, we evidence for the first indication of a Kerr BH or, possibly, a Kerr-Newman BH formation: \begin{itemize} \item The P-GRB spectrum of GRB 090510, in the time interval from $T_0+0.528$ to $T_0+0.644$ s, is best-fitted by a Comptonized component (see figures~\ref{PGRB} and \ref{comp} and table~\ref{tab:fit}), which is interpreted as a convolution of thermal spectra originating in a dyadotorus (see \citealt{Cherubini2009} and \citealt{Ruffini2009}, figure~\ref{dyadotorus}, and section 2). \item The prompt emission follows at the end of the P-GRB (see figure~\ref{spectotal}). The analysis of the prompt emission within the fireshell model allows to determine the inhomogeneities in the CBM giving rise to the spiky structure of the prompt emission and to estimate as well an averaged CBM density of $\langle n_{\rm CBM}\rangle=8.7\times10^{-6}$~cm$^{-3}$ obtained from a few CBM clouds of mass $\sim10^{22}$ g and typical dimensions of $\sim10^{16}$ cm (see figure~\ref{090510simlc}). Such a density is typical of galactic halos where binary NS are expected to migrate due to large natal kicks. \item The late X-ray emission of GRB 090510 does not follow the characteristic patterns expected in BdHN events (see figure~\ref{episode3} and \citealt{Pisani2013}). \item The GeV emission occurs at the end of the P-GRB emission and is initially concurrent with the prompt emission. This sequence occurs in both S-GRBs \citep{Muccino2015} and BdHNe \citep{Wang2015}. This delayed long lasting ($\approx200$ s) GeV emission in GRB 090510 is one of the most intense ever observed in any GRB \citep[see figure~\ref{090510GeV} and][]{Ackermann2013,Ruffini2016}. \item We then consider accretion on co-rotating and counter-rotating orbits (see Ruffini \& Wheeler 1969, in problem 2 of $\S$ 104 in \citealt{LL2003}) around an extreme Kerr BH. Assuming the accretion of the crustal mass $2M_c=8.60\times10^{-5}$~M$_\odot$ from a $1.6+1.6$~M$_\odot$ NS--NS binary, fulfilling global charge neutrality (see figure~\ref{Nanda}), geometrical beaming angles of $\theta\gtrsim0^{\rm o}.81$, for co-rotating case, and $\theta\gtrsim2^{\rm o}.70$, for the counter-rotating case, are inferred. In order to fulfill the transparency condition, the initial Lorentz factor of the jetted material has to be $\gamma\gtrsim550$ (see section 6.6). \item While there is evidence that the GeV emission must be jetted, no beaming appears to be present in the P-GRB and in the prompt emission, with important consequence for the estimate of the rate of such events \citep{Ruffini2016}. \item The energetic and the possible beaming of the GeV emission requires the presence of a Kerr BH, or a Kerr-Newman BH dominated by its angular momentum and with electromagnetic fields not influencing the geometry (see also section 6.5). \item The self-consistency of the entire procedure has been verified by estimating, on the ground of the fireshell theory, the cosmological redshift of the source. The theoretical redshift is $z=0.75\pm0.17$ (see section~6.4), close to and consistent with the spectroscopically measured value $z=0.903\pm0.003$ \citep{Rau2009}. \item The values of $E_{\mathrm{peak}}$ and $E_{\mathrm{iso}}$ of GRB 090510 fulfill with excellent agreement the MuRuWaZha relation \citep[see section~5.2, figure~\ref{calderone} and][]{RuffiniGCN}. \end{itemize} The main result of this article is that the dyadotorus manifests itself by the P-GRB emission and clearly preceeds the prompt emission phase, as well as the GeV emission originating from the newly-formed BH. This contrasts with the usual assumption made in almost the totality of works relating BHs and GRBs in which the BH preceeds the GRB emission. In conclusion, in this article, we take GRB 090510 as the prototype of S-GRBs and perform a new time-resoved spectral analysis, in excellent agreement with that performed by the AGILE and the \textit{Fermi} teams. Now this analysis, guided by a theoretical approach successfully tested in this new family of S-GRBs, is directed to identify a precise sequence of different events made possible by the exceptional quality of the data of GRB 090510. This include a new structure in the thermal emission of the P-GRB emission, followed by the onset of the GeV emission linked to the BH formation, allowing, as well, to derive the strucutre of the circumburst medium from the spiky structure of the prompt emission. This sequence, for the first time, illustrates the formation process of a BH. It is expected that this very unique condition of generating a jetted GeV emission in such a well defined scenario of a newly-born BH will possibly lead to a deeper understanding of the equally jetted GeV emission observed, but not yet explained, in a variety of systems harboring a Kerr BH. Among these systems we recall binary X-ray sources \citep[see, e.g.,][and references therein]{1978pans.proc.....G}, microquasars \citep[see, e.g.,][and references therein]{2015IAUS..313..370M}, as well as, at larger scale, active galactic nuclei \citep[see e.g.,][and references therein]{2015A&A...579A..34A}.

1607.02400

1607

1607.07822_arXiv.txt

The {\it Fermi} Large Area Telescope (LAT) is currently the most important facility for investigating the GeV $\gamma$-ray sky. With {\it Fermi} LAT more than three thousand $\gamma$-ray sources have been discovered so far. 1144 ($\sim40\%$) of the sources are active galaxies of the blazar class, and 573 ($\sim20\%$) are listed as Blazar Candidate of Uncertain type (BCU), or sources without a conclusive classification. We use the Empirical Cumulative Distribution Functions (ECDF) and the Artificial Neural Networks (ANN) for a fast method of screening and classification for BCUs based on data collected at $\gamma$-ray energies only, when rigorous multiwavelength analysis is not available. Based on our method, we classify 342 BCUs as BL Lacs and 154 as FSRQs, while 77 objects remain uncertain. Moreover, radio analysis and direct observations in ground-based optical observatories are used as counterparts to the statistical classifications to validate the method. This approach is of interest because of the increasing number of unclassified sources in {\it Fermi} catalogs and because blazars and in particular their subclass High Synchrotron Peak (HSP) objects are the main targets of atmospheric Cherenkov telescopes.

Blazars are active galactic nuclei (AGN) with a radio-loud behavior and a relativistic jet pointing toward the observer. \citep{Abdo01} These sources are divided into two main classes: BL Lacertae objects (BL Lacs) and Flat Spectrum Radio Quasars (FSRQs), which show very different optical spectra even if in other wavebands they are similar. FSRQs have strong, broad emission lines at optical wavelengths, while BL Lacs show at most weak emission lines, sometimes display absorption features, and can also be completely featureless. Compact radio cores, flat radio spectra, high brightness temperatures, superluminal motion, high polarization, and strong and rapid variability are commonly found in both BL Lacs and FSRQs. Blazars emit variable, non-thermal radiation across the whole electromagnetic spectrum, which includes two components forming two broad humps in a $\nu f{_\nu}$ representation. The low-energy one is attributed to synchrotron radiation, and the high-energy one is usually thought to be due to inverse Compton radiation. See \citet{Ghisellini} for a recent review of the properties of $\gamma$-ray AGN. Blazars can also be classified into different subclasses based on the position of the peak of the synchrotron bump in their spectral energy distribution (SED), namely, low frequency peaked (LSP or sources with $\nu^{S}_{peak}$ < $10^{14}$ Hz), intermediate frequency peaked (ISP or sources with $10^{14}$ Hz < $\nu^{S}_{peak}$ < $10^{15}$ Hz) and high frequency peaked (HSP or sources with $\nu^{S}_{peak}$ > $10^{15}$ Hz ) \citep{Abdo02}. This subclassification suggests the possibility that the $\gamma$-ray properties of the sources may lead to constraints on the type of objects responsible for the radiation especially in view of the increasing number of detections obtained by the {\it Fermi} Large Area Telescope (LAT) that still have to be properly classified. The Third {\it Fermi}-LAT Source Catalog (3FGL)\citep{ace15} listed 3033 $\gamma$-ray sources collected in four years of operation, from 2008 August 4 (MJD 54682) to 2012 July 31 (MJD 56139). 3FGL covers the full sky. 1144 sources are identified or associated with galaxies of the blazar class. 660 are BL Lacs and 484 are FSRQs. 3FGL includes also 573 Blazar Candidates of Uncertain type (BCUs). Because of the difficulty of having extensive optical observation campaigns for full classification of blazars, if we compare the 3FGL with previous catalogs released by the LAT collaboration we can see a significant increase of the number of unclassified blazars. In Table 1 we show the growth of the number of blazar-class sources in {\it Fermi}-LAT catalogs and the relative fraction of each blazar source subclass. The percentage of BCUs within the blazar sample increased from 13.8$\%$ in 1FGL to 33.4$\%$ in 3FGL. Although the detailed multiwavelength analysis necessary for unambiguous classification has been done and is continuing for many of these \citep{alvarez}, a first classifying screening of BCUs, as our method proposes, can be very useful for the blazar scientific community. \begin{table} \caption{Blazar-class source distribution in {\it Fermi}-LAT catalogs and the relative fraction of each blazar source subclass} \begin{center} \begin{footnotesize} \begin{tabular}{lccr} \hline \hline \bf{Class} &\bf{1FGL} & \bf{2FGL} & \bf{3FGL} \\ \hline BLL & 295 (44.4\%) & 436 (41\%) & 660 (38.4\%)\\ FSRQ & 278 (41.8\%) & 370 (34.8\%) & 484 (28.2\%)\\ BCU & 92 (13.8\%) & 257 (24.2\%) & 573 (33.4\%)\\ \hline Total & 665 & 1063 & 1717\\ \hline \end{tabular} \end{footnotesize} \end{center} \end{table} The aim of this work is to find a simple estimator in order to classify BCUs and, when it is possible, to identify high-confidence HSP candidates. The present generation of Imaging Atmospheric Cherenkov Telescopes (IACTs), such as VERITAS, H.E.S.S. and MAGIC, has opened the realm of ground-based $\gamma$-ray astronomy in the Very High Energy range (VHE: E $> $100 GeV). The Cherenkov Telescope Array (CTA) will explore our Universe in depth in this energy band and lower. For a recent review of present and future Cherenkov telescopes, see \citep{deNaurois}. The BL Lac HSP sources are the most numerous class of TeV sources. The TeV catalog \citep{horan} reports 176 TeV sources. 46 of them are HSP BL Lacs and only 5 FSRQs, therefore the ability to correctly identify HSP objects will be very important for the Cherenkov scientific community and in the determination of CTA targets, in order to increase the rate of detections, since IACTs have a small field of view.The novelty of the present approach is that our study relies exclusively on variability data collected at $\gamma$-ray energies where {\it Fermi}-LAT is most sensitive (0.1 -- 100 GeV) and it remains totally independent from other data at different wavelengths. The paper is laid out as follows: in Section \ref{sec:LAT}, we present the $\gamma$-ray data and the ECDF light curves considered for our analysis; in Sect.~\ref{sec:ANN}, we describe the use of artificial neural networks, and in Sect.~\ref{sec:results} we present the results of the ANN analysis. In Sect.~\ref{sec:class} we present a summary of the results of our classification of BCUs listed in the 3FGL {\it Fermi}-LAT and we highlight the most promising HSP candidates. In Sect.~\ref{sec:multiwave} we test our method comparing the predicted classifications with additional data, obtained through optical spectroscopy and radio observations. We summarize our conclusions in Sect.~\ref{sec:conclusions}. \section {Gamma-ray Data} \label{sec:LAT} \subsection{The Large Area Telescope} The LAT is the primary instrument on the {\it Fermi Gamma-ray Space Telescope}, launched by NASA on 2008 June 11 and it is the first imaging {\bf GeV $\gamma$-ray observatory} able to survey the entire sky every day at high sensitivity orbiting the Earth every 96 minutes. The \emph{Fermi} LAT is a pair-conversion telescope with a precision converter-tracker and calorimeter. It measures the tracks of the electron and positron that result when an incident $\gamma$ ray undergoes pair-conversion and measures the energy of the subsequent electromagnetic shower that develops in the telescope's calorimeter \citep{FLAT}. Data obtained with \emph{Fermi}-LAT permit rapid notification and facilitate monitoring of variable sources such as the BCUs that we consider in this study. In this paper we used the monthly $\gamma$ flux value from LAT 4-year Point Source Catalog (3FGL) and the Fermi Science Support Center (FSSC) for any other data\footnote{\tt http://fermi.gsfc.nasa.gov/ssc/data/access/lat/4yr\_catalog/}. \subsection{B-FlaP: Blazar Flaring Patterns } Variability is one of the defining characteristics of blazars \citep{Paggi01}. We considered the light curves of the blazar sources evaluated with monthly binning, as reported in 3FGL catalog, and with these data we designed the basic structure of the B-FlaP method.\\ The original idea was to compare the $\gamma$-ray light curve of the source under investigation with a \emph{template} classified blazar class light curve, then measure the difference in a proper metric. Typically $\gamma$-ray AGN are characterized by fast {\bf flaring} that could alter significantly the light curve and could make the comparison difficult. In addition, different flux levels could hide the actual similarity of light curves. As first approach of this study we compute the Empirical Cumulative Distribution Function (ECDF) of the light curves \citep{KS}. We constructed the percentage of time when a source was below a given flux by sorting the data in ascending order of flux and then compared the ECDFs of BCUs with the ECDFs of blazars whose class is already established, (\S\ref{sec:LAT}). This is our variation of the Empirical Cumulative Distribution Function (ECDF) method. In Fig.~\ref{<Fig.1>} we show the ECDF plots for 3FGL blazars and BCUs. In principle, differences due to the flaring patterns of BL Lacs and FSRQ appear in two ways: (1) the flux where the percentage reaches 100 represents the brightest flare seen for the source; and (2) the shape of the cumulative distribution curve reveals information about the flaring pattern, whether the source had one large flare, multiple flares, or few flares.The BL Lacs have fewer large flares than the FSRQs, and the FSRQ curves are more jagged, suggesting multiple flares compared to the smoother BL Lac curves. The difference between the classes is observed when we plotted the two blazar classes together. At the bottom left of Fig.~\ref{<Fig.1>} is shown the significant overlap between the types where it is hard to distinguish individual objects, and there are outliers that extend beyond the range of the plots, but it is possible to recognize on the top left of the diagram a specific area where the overlap between BL Lac and FSRQ is minimal. This area, at values of the flux less than $\sim$ 2.5 $\times$ 10$^{-8}$ ph cm$^{-2}$ s$^{-1}$, could lead to a first \emph{qualitative} recognition of BL Lac objects. In B-FlaP, special attention is needed for upper limits, which arise whenever light curves are constructed with fixed binning, as is the case here. They can be naturally incorporated into the current ECDF method, as the points plotted in the diagrams are the percentage of time that the source is below a given flux value. Nevertheless, upper limits could introduce biases, skewing the cumulative distribution toward higher percentages. Upper limits could be avoided entirely by producing light curves with adaptive binning \citep{Benoit02}, a technique that could be implemented into a possible follow-up study. For this reason and because the ECDF plots represent only a proof of concept of the whole method, we follow up the ECDF first analysis with an Artificial Neural Network analysis (ANN) by an original algorithm developed to distinguish the single BCU object and to give its likelihood to be a BL Lac or a FSRQ. \\ The reasons for the flaring patterns differences between BL Lacs and FSRQ are very likely connected with the processes occurring at the base of the jet, where the largest concentration of relativistic particles and energetic seed photons are expected. While in FSRQs accretion onto the central black hole produces a prominent and variable spectrum, characterized by continuum and emission-line photons, usually accompanied by the ejection of relativistic blobs of plasma in the jet, BL Lacs do not show such kind of activity and most of the observed radiation originates within the jet itself. As a consequence, the production of $\gamma$-ray emission through inverse Compton (IC) scattering can change much more dramatically in FSRQs than in BL Lac-type sources, where the contribution of the central engine to the seed radiation field is weaker \citep{VAR} \begin{figure*} \begin{center} {\includegraphics[width=0.8\textwidth]{ECDF_DAVID.png}} \caption{ECDF plots of {\it Fermi} Blazars: BL Lacs (top left), FSRQs (top right) , BL Lac and FSRQ overlap (bottom left), BCUs (bottom right). The cumulative percentage of bins with flux below a given level is shown as a function of the 0.1 -- 100 GeV flux in a bin, in units of 10$^{-8}$ ph cm$^{-2}$ s$^{-1}$.\label{<Fig.1>}} \end{center} \end{figure*} \subsubsection{High Synchrotron Peak blazar} With reference to the aim of this study we applied the same ECDF technique to the blazar subclasses. Using the Third Catalog of Active Galactic Nuclei detected by Fermi-LAT \citep[3LAC,][]{Ackermann2015}, we collected information about classification and SED distribution of the blazars. The third release of the catalog considers only 1591 AGN detected at |\emph{b}| >$\ang{10}$ where \emph{b} is the Galactic latitude, 289 are classified sources as HSP on the basis of their SED, where 286 of them are represented by BL Lac objects and 3 by FSRQs. 160 of the 573 BCUs are HSP suspects. For all the other data in this study we referred to 3FGL. While ISP and LSP blazars show the most variable patterns and can belong to both the BL Lac or FSRQ families, HSP objects are characterized by nearly constant emission. In Fig.~2 we plotted the ECDF for 3LAC HSPs versus FSRQs. As we expected, because of the fact that HSPs are almost exclusively represented by BL Lac objects, the HSPs went through the BLL clean area at the upper left corner of the plot. Even if ISP and LSP contamination is not negligible (Fig.~3), the result observed in Fig.~2 suggests \emph{the potential ability} of ECDF B-FlaP to identify a flux range at the 100th percentile (less than $\sim$ 2.0 $\times$ 10$^{-8}$ ph cm$^{-2}$ s$^{-1}$) where it is possible to not only determine the blazar class but also to tentatively assign the HSP subclass for a BCU source. However, even here, visual inspection of the curves in all the ECDF figures suggests that the shape of the curve does not show major differences between the observed blazar classes. In order to improve the analysis we used the same ANN algorithm developed for BCUs for the HSP classification. \begin{figure} \resizebox{\hsize}{!}{\includegraphics{hspFSRQ3fgl_DAVID.png}} \caption{ECDF for HSPs (black) and FSRQs (red), using the same construction and scale as Fig.1} \end{figure} \begin{figure} \resizebox{\hsize}{!}{\includegraphics{HSPlspISP.png}} \caption{LSPs and ISPs (red) versus HSPs (blue). Flux distribution at the 100th percentile. Even if, in the considered $\gamma$-ray flux range the number of HSPs is greater than the other subclasses, the degree of contamination cannot be overlooked.} \end{figure}

\label{sec:conclusions} We developed an Artificial Neural Network technique based on B-FlaP information to assess the likelihood that a $\gamma$-ray source can belong to the BL Lac or FSRQ family, or more interestingly to the HSP blazar subclass, using only $\gamma$-ray data. We tested the method on sources that are classified as BL Lacs or FSRQs in 3FGL, and focusing our attention on the HSP blazars, we found, and here we confirm, that the method is very effective in the identification of blazars and offers an opportunity to provide a tentative HSP classification. This paper presents a full classification list of the 3FGL BCUs according to their ANN likelihood. BCUs can be divided in 342 BL Lac candidates (\bann), 154 FSRQ candidates (\fann) leaving only 77 without a clear classification. Among 573 BCUs we selected and ranked 53 very interesting HSP candidates to be observed through Very High Energy Telescope. In order to validate the method we compared B-FlaP with the Variability Index and the Power Law Index. In both the comparisons B-FlaP showed full consistency, and in some cases the efficiency of B-FlaP is greater than what is obtainable by the other parameters. To further assess the reliability of the method we performed direct optical observations for a sample of BCUs with Galactic latitude $|b| > 10^{\rm o}$ and maximum $\gamma$-ray flux less than $6 \cdot 10^{-8}\, {\rm ph\, cm^{-2}\, s^{-1}}$ . In those cases where we were able to perform spectroscopic observations we found that the optical spectra were fully consistent with the expectations based on the ANN results. Even the results of benchmarking between the radio data and B-FlaP showed a consistency of assessment with the two approaches. We conclude that, although B-FlaP cannot replace confirmed and rigorous spectroscopic techniques for blazar classification, it may be configured as an additional powerful approach for the preliminary and reliable identification of BCUs and in particular the HSP blazar subclass when detailed observational or multiwavelength data are not yet available.

1607.07822

1607

1607.06129_arXiv.txt

Braking indices of pulsars present a scientific challenge as their theoretical calculation is still an open problem. In this paper we report results of a study regarding such calculation which adapts the canonical model (which admits that pulsars are rotating magnetic dipoles) basically by introducing a compensating component in the energy conservation equation of the system. This component would correspond to an effective force that varies with the first power of the tangential velocity of the pulsar's crust. We test the proposed model using data available and predict braking indices values for different stars. We comment on the high braking index recently measured of the pulsar J1640-4631.

Pulsars are normally modelled as rapidly rotating, highly magnetized stars composed mainly of neutrons. It has been observed that their rotation frequencies are decaying, this spin-down being quantified by the braking index (BI), $n$, defined by: \eqnn{eq:def_n}{n \equiv \frac{{\Omega \ddot \Omega }} {{\dot \Omega ^2 }},} where $\Omega$ is the pulsar's angular velocity and the dot denotes a time derivative. In such model, which we will refer to as canonical, the main time-varying field responsible for the loss of rotational energy in a pulsar is a magnetic dipole field \citep{osg}. Also, the canonical model predicts n=3 for all existing pulsars. There are not many pulsars for which the BI was obtained observationally (see Table \ref{tab:pulsars}). Most of their BI values lie within the range $0.9 - 2.8$ (see Table \ref{tab:bComplete}). The only pulsar with index greater than three is J1640-4631, whose value, $n=3.15$, was recently measured\citep{ar16}. Since the canonical model fails to yield the observed BI, improvements on this model have been tried involving different theoretical approaches \citep{br88, ah97, me97, cs06, mag12, kt15}. \begin{table*} \centering \caption{ Angular velocities ($\Omega$) for pulsars with known braking indices. Time derivatives are denoted by a dot. } \label{tab:pulsars} \begin{tabular}{cccccc} \hline PSR & $\Omega$ & $\dot{\Omega}$ ($\times 10^{-10}$ &$\ddot{\Omega}$ ($\times 10^{-21}$ & References \\ & (rad s$^{-1}$)& rad s$^{-2}$) & rad s$^{-3}$) & & \\ \hline B0531+21 (Crab) & 189.912022 & -24.2674 &78.075 & \citet{lyne93,lyne15} \\ B0833-45 (Vela) & 70.4 & -0.986 & 0.19 & \citet{lyne96,lyne15} \\ B0540-69 & 124.623817 & -11.8365 & 24.1 & \citet{liv2007} \\ B1509-58 & 41.68013054 & -4.24618765 & 12.2944 &\citet{liv2007} \\ J1846-0258 & 19.340994108 & -4.21955 & 24.3 & \citet{liv2007} \\ J1119-6127 & 15.401361301 & -1.517708 & 4.014 & \citet{welte2011} \\ J1734-3333 & 5.37327178 & -0.104742 & 0.018 & \citet{espinoza2011} \\ J1833-1034 & 101.5322352 & -3.314451250 & 2.008734343 & \citet{rgl12}\\ J1640-4631 & 30.4320477075 & -1.433053 & 2.12 & \citet{ar16}\\ \hline \end{tabular} \end{table*} In this paper we analyse a modification of the canonical model aiming at theoretically obtaining BI of pulsars that have already been observed. The modification consists basically on introducing a compensating component in the energy conservation equation of the system. This component would correspond to a force that varies with the first power of the tangential velocity of the pulsar's crust. In the next section we present a summary of the canonical model followed by a description of its modified version as proposed by us. The remaining sections provide the results using the modified model and their analysis. We close the paper with our concluding remarks.

We analysed a modification of the canonical model for pulsars' spin-down that introduces an effective force that is tangential to the star's rotation motion. Our goal was to provide more precise predictions of pulsars' BI. The effective force involves all yet unknown physical contributions that make pulsars' BI less than three. Our results were possible assuming that the pulsars have the same (typical) values for some physical characteristics, like mass and radius. We used eight pulsars with observed BI to calibrate the model. Also, we discovered an extra relation between the BI, the ratio $|\dot \Omega|/\Omega$ and the tangential force constant of a pulsar which enabled us to make predictions of BI of other pulsars. The results that we found are applicable to pulsars with $|\dot \Omega|$ larger than 10$^{-11}$ rad s$^{-2}$ and with $|\dot \Omega|/\Omega$ near the range 1 - 25 $\times$ 10$^{-12}$ s$^{-1}$. In order to improve the model, it is important to find physical details about the effective force. Its mathematical structure is simple and general, allowing different physical possibilities for its origin. By using in our model data of the high-n pulsar J1640-4631 we found evidence that the model can also indicate pulsars with BI larger than three. $ \\$ N.S.M., A.S.O. and C.F. acknowledge the Brazilian federal funding agency CNPq for financial support (grants 309295/2009-2, 149107/2010-2 and 312906/2013-7, respectively) as well as the National Institute of Science and Technology in Astrophysics (INCT-A, Brazil) and FAPESP (thematic project, grant 13/26258-4).

1607.06129

1607

1607.01824_arXiv.txt

We present a theoretical model aimed at explaining the IRX-$\beta$ relation for high redshift ($z \simgt 5$) galaxies. Recent observations \citep{Capak15, Bouwens16} have shown that early Lyman Break Galaxies, although characterized by a large UV attenuation (e.g. flat UV $\beta$ slopes), show a striking FIR deficit, i.e. they are ``infrared-dark''. This marked deviation from the local IRX-$\beta$ relation can be explained by the larger molecular gas content of these systems. While dust in the diffuse ISM attains relatively high temperatures ($T_d \simeq 45$ K for typical size $a=0.1 \mu$m; smaller grains can reach $T_d \simgt 60 $ K), a sizable fraction of the dust mass is embedded in dense gas, and therefore remains cold. If confirmed, the FIR deficit might represent a novel, powerful \textit{indicator of the molecular content of high-$z$ galaxies} which can be used to pre-select candidates for follow-up deep CO observations. Thus, high-$z$ CO line searches with ALMA might be much more promising than currently thought.

\label{Mot} Dust grains are a fundamental constituent of the interstellar medium (ISM) of galaxies. A large fraction ($\approx $ 50\% in the Milky Way) of the heavy elements produced by nucleosynthetic processes in stellar interiors can be locked into these solid particles. They are vital elements of the ISM multiphase gas life-cycle and key species for star formation, as they catalyze the formation of molecules on their surfaces. Most relevant here, they efficiently absorb optical/ultraviolet (UV) stellar light, by which they are heated, and re-emit this energy as longer (far-infrared, FIR) wavelength radiation that can freely escape from the galaxy. It is then natural to expect a tight relation between the UV ``deficit'' and the IR excess produced by this process. Indeed such correlation, called the IRX-$\beta$ relation \citep{Meurer95, Meurer99, Calzetti00, Takeuchi12, Grasha13}, links the UV spectral slope (defined as $\propto \lambda^\beta$) with the FIR excess, i.e. the ratio between the total FIR and $\approx 1500$ \AA\ UV fluxes, $F_{FIR}/F_{UV}$. Such relation has been extensively applied to the local Universe, and more recently extended to redshift $z \approx 2$ \citep{Reddy12, Alvarez16, Fujimoto16}. The relation has proven to be very robust, at least for starburst galaxies, and useful as it allows to estimate the total UV attenuation. Conversely, if the intrinsic and observed $\beta$ values are known, one can derive the expected total FIR luminosity. The presence of dust at high ($z\simgt 6$) redshift implies that conventional dust sources (AGB and evolved stars) are not the dominant contributors. This is because their evolutionary timescales are close or exceed the Hubble time at that epoch ($\approx 1$ Gyr). Following the original proposal by \cite{Todini01}, it is now believed that the first cosmic dust were formed in the supernova ejecta ending the evolution of much more fast-evolving massive stars \citep{Hirashita02, Nozawa07, Bianchi07, Gall11}. For similar reasons the standard grain growth suffered by grains during their residence time in molecular clouds (MC) of contemporary galaxies cannot increase the amount of dust by considerable amounts \citep{Ferrara16}. Thus, albeit quasar host galaxies show remarkably high dust masses \citep{Beelen06, Michalowski10}, in general the dust-to-gas ratio towards high-redshift rapidly decreases \citep{Dunlop13} as also witnessed by the observed steepening of early galaxies UV spectra. This does not come as a complete surprise given that the average metallicity of the Universe increases with time. \cite{Ferrara99} (for a recent calculation see \cite{DaCunha13}) noticed another important feature of high-$z$ dust. Due to the redshift increase of CMB temperature, $\Tcmb= T_0 (1+z)$ K with $T_0=2.725$, the FIR signal from dust becomes increasingly swamped by the CMB. At $z=6$, for example, $\Tcmb=19.07$ K; as usually dust temperatures in the \textit{diffuse} ISM of galaxies are in the range $20-40$ K, the effect cannot be neglected. Even more dramatic, if not complete, might be the suppression of the signal from dust in dense regions (e.g. molecular clouds) where the dust is in thermal equilibrium with the CMB. In the light of these physical ideas, here we intend to revisit the interpretation of recent dust detections at redshift $z\simgt 5$. The superb sensitivity of the ALMA interferometry has allowed to detect the FIR signal of a handful of Lyman Break Galaxies (LBGs) for which HST rest-frame UV photometry (and hence a $\beta$ determination) is available \citep{Capak15}. This experiment has reported a puzzling deviation of detected LBGs from the more local IRX-$\beta$ relation. In practice, these galaxies, although characterized by relatively flat $\beta \approx -1$ values, indicative of non-negligible dust attenuation, show a noticeable FIR deficit, i.e. they are relatively ``infrared-dark''. Such suggested deficit has been strongly reinforced by an even more recent report by the ASPECS survey \citep{Bouwens16}. The authors have performed deep 1.2 mm-continuum observations of the Hubble Ultra Deep Field (HUDF) to probe dust-enshrouded star formation from 330 LBGs spanning the redshift range $z = 2-10$. The striking result is that the expectation from the Meurer IRX-$\beta$ relation at $z=4$ was to detect at least 35 galaxies. Instead, the experiment only provided 6 tentative detections (in the most massive galaxies of the sample). Clearly, redshift evolution either of the dust temperature and/or mass must play a key role. Understanding this behavior is the central aim of the present study\footnote{Throughout the paper, we assume a flat Universe with the following cosmological parameters: $\Omega_{\rm m} = 0.308$, $\Omega_{\Lambda} = 1- \Omega_{\rm m} = 0.692$, and $\Omega_{\rm b} = 0.048$, where $\Omega_{\rm M}$, $\Omega_{\Lambda}$, $\Omega_{\rm b}$ are the total matter, vacuum, and baryonic densities, in units of the critical density, and $h$ is the Hubble constant in units of 100 km/s \citep{Ade15}. }. An exception to the above scenario is the puzzling case of A1689-zD1 \citep{Watson15, Knudsen16}, a $z=7.5\pm 0.2$ gravitationally-lensed LBG for which the thermal dust emission has been detected by ALMA. The large FIR flux $L_\mathrm{FIR}=(6.2\pm 0.8)\times 10^{10}$ L$_{\sun}$ is indicative of considerable amounts of dust, consistent with a Milky Way dust-to-gas ratio. The rest-frame UV slope is estimated to be $\beta=-2\pm 0.1$. In this paper we propose a novel idea to explain the physics behind the IRX-$\beta$ relation evolution. Moreover, we will introduce a new indicator which can be used to infer the molecular mass of early galaxies in a regime where classical tracers as CO emission lines might be unavailable, difficult to obtain, or affected by increasing uncertainties (as, e.g., the CO-to-H$_2$ conversion factor).

\label{Con} We have presented a dust extinction and FIR emission model that is aimed at explaining the IRX-$\beta$ relation for high redshift ($z \simgt 5$) galaxies. We have first derived the dust mass vs. UV spectral slope, $\beta$ relation. Then the temperature of grains exposed to the typical range of internal UV field intensities of $z\simeq 6$ LBG galaxies has been computed. Such calculation allows for attenuation of the UV flux inside molecular clouds and account for the CMB effects both for what concerns dust temperature and dust continuum suppression. When combined, these two results allow to predict the expected IRX-$\beta$ relation (assuming a SMC extinction curve) and compare it with recent data. A key result is that our model reproduces extremely well (essentially without free parameters) the striking FIR deficit observed, showing that these early systems are ``infrared-dark''. We suggest that the deficit is caused by an increasing molecular gas content of these systems. While dust residing in the diffuse ISM attains large temperatures ($T_d \simeq 45$ K for typical size $a=0.1 \mu$m; smaller grains can reach $T_d \simgt 60 $ K) due to presence of intense interstellar UV fields \citep{Bethermin15}, dust located in molecular clouds becomes very cold (but not colder than the CMB). As galaxies with a larger molecular content should be characterized by a lower IRX, our model suggests an interesting way to pre-select candidates for molecular studies at high-$z$ studies. In other words, for a given value of $\beta$, we predict that the galaxies with the largest molecular content are those characterized by the lower IRX values. That is, IRX anti-correlates with the molecular fraction, $\mu$. Thus, searching for CO line emission (and perhaps H$_2$ lines directly with, e.g., SPICA\footnote{www.ir.isas.jaxa.jp/SPICA/}) from high-$z$ galaxies with ALMA might be much more promising than currently thought. A first attempt along these lines has been presented by \cite{Riechers14} who observed the LBG galaxy HZ6 (IRX = 0.38, $\beta=-1.13$), also part of the Capak sample. The authors observed HZ6 with the VLA in search for CO($J = 2-1$) line emission, obtaining an approximate 3$\sigma$ limit of $<0.03$ Jy km s$^{-1}$. However, higher-$J$ transitions, like the CO(6-5), detectable with ALMA, are expected to be considerably more luminous \citep{Vallini16}. If successful these experiments could provide a fundamental insight on the highly debated processes of dust formation and growth \citep{Ferrara16}, on the evolution of the molecular content of galaxies \citep{Genzel15} and the relationship between the two.

1607.01824

1607

1607.08591_arXiv.txt

We use numerical simulations to measure the sensitivity of the tidal spin down rate of a homogeneous triaxial ellipsoid to its axis ratios by comparing the drift rate in orbital semi-major axis to that of a spherical body with the same mass, volume and simulated rheology. We use a mass-spring model approximating a viscoelastic body spinning around its shortest body axis, with spin aligned with orbital spin axis, and in circular orbit about a point mass. The torque or drift rate can be estimated from that predicted for a sphere with equivalent volume if multiplied by $0.5 (1 + b^4/a^4)(b/a)^{-4/3} (c/a)^{-\alpha_c}$ where $b/a$ and $c/a$ are the body axis ratios and index $\alpha_c \approx 1.05$ is consistent with the random lattice mass spring model simulations but $\alpha_c = 4/3$ suggested by scaling estimates. A homogeneous body with axis ratios 0.5 and and 0.8, like Haumea, has orbital semi-major axis drift rate about twice as fast as a spherical body with the same mass, volume and material properties. A simulation approximating a mostly rocky body but with 20\% of its mass as ice concentrated at its ends has a drift rate 10 times faster than the equivalent homogeneous rocky sphere. However, this increase in drift rate is not enough to allow Haumea's satellite, Hi'iaka, to have tidally drifted away from Haumea to its current orbital semi-major axis.

The simplicity of analytic formulae for tidal spin down of spherical bodies and the associated drift rate in semi-major axis \citep{goldreich63,goldreich68,kaula64,efroimsky09,efroimsky13} has made it possible to estimate tidal spin-down rates in diverse settings, including stars, exoplanets, satellites, and asteroids (e.g., \citealt{yoder81,efrolainey,ferrazmello08,ogilvie14}). However, there is uncertainty in how to estimate the spin down rate and associated semi-major axis drift of non-spherical bodies (though see \citealt{mathis09} for estimates of tidal drift rates for two extended homogeneous bodies) and this hampers attempts to interpret the dynamical history of systems that include extended spinning bodies like Pluto's and Haumea's satellite systems (e.g., \citealt{ragozzine09,cuk13,weaver16}). An extended shape might experience enhanced tidal distortion and dissipation compared to a stronger spherical body. \citet{ragozzine09} suggested that use of the radius of an object with equivalent volume (the volumetric radius) in classic tidal formulae could lead to an underestimate of the tidal torque on Haumea. A semi-analytical treatment of the gravitational potential inside a triaxial homogeneous body can give a description of instantaneous internal deformations (or displacements) and stresses as a function of Cartesian coordinates \citep{dobrovolskis82}. If the axes of the tidal force lie along planes of body symmetry, the displacements depend on 12 dimensionless coefficients that can be computed numerically. However, analytical computation in the general case is more challenging. ``The reader is cautioned that in the general case where the tide-raising object does not lie along a principal axis, as many as 72 unknown coefficients may appear.'' \citep{dobrovolskis82}. To semi-analytically compute the tidal torque for a spinning triaxial body we would need to numerically compute all these coefficients and then average over body rotation to compute secular or average drift rates. Due to the complexity of accurate semi-analytical computation, it would be convenient to have scaling relations, dependent on body axis ratios, that would allow one to obtain correct analytical formulae for tidal evolution from the well known one for homogeneous spherical bodies. Our goal here is to numerically measure such correction factors using numerical simulations. Because of their simplicity and speed, compared to more computationally intensive grid-based or finite element methods, mass-spring computations are an attractive method for simulating deformable bodies. By including spring damping forces, they can be used to model viscoelastic deformation. We previously used a mass-spring model to study tidal encounters \citep{quillen16} and measure tidal spin down for spherical bodies over a range of viscoelastic relaxation timescales \citep{frouard16}. Mass spring models are not restricted to spherical particle distributions and so can be used to study triaxial ellipsoids. Our work \citep{frouard16} compared orbital semi-major axis drift rates for spherical bodies to those predicted analytically. Here we use the same type of simulations to measure drift rates but work in a setting where analytical computations are lacking and the numerical measurements may motivate order of magnitude scaling arguments.

\begin{table} \vbox to90mm{\vfil \caption{\large Information about Haumea and Hi'iaka \label{tab:haumea}} \begin{tabular}{@{}lllllll} \hline Haumea: && \\ Semi-major axis of ellipsoid & $a_H$ & 960 km \\ Axis ratio & $b_H/a_H$ & 0.80 \\ Axis ratio & $c_H/a_H$ & 0.52 \\ Volumetric radius & $R_{vH}$ & 716.6 km \\ Mass of Haumea & $m_H$ & $4 \times 10^{21}$ kg\\ Energy density scale & $e_{g,H}$ & 4.05 GPa \\ Gravitational timescale & $t_g$ & 1171 s \\ Spin rate & $\sigma_{H} t_g$ & 0.52 \\ Tidal frequency & $\omega \sim 2 \sigma_{H}$ & $0.9 \times 10^{-3}$ Hz\\ \hline Mass ratio & $q = M_{Hi}/M_{H}$ & 0.0045 \\ orbital semi-major axis Hi'iaka & $a_{Hi}$ & 49880 km \\ \hline \end{tabular} {\\ Body semi-major axis and axis ratios are by \citet{lockwood14}. Mass of Haumea, mass ratio, $q$, of Hi'iaka and Haumea and semi-major axis are by \citet{ragozzine09}. The volumetric radius $R_{vH}$ is the radius of a sphere with the same volume as the triaxial ellipsoid. Spin rate, gravitational timescale, $t_g$ and energy density scale, $e_g$, are computed using the volumetric radius and equations \ref{eqn:tgrav} and \ref{eqn:eg}. The spin rate was computed using the spin period $ P_H = 2\pi/\sigma_H= 3.91531 \pm 0.00005$ hours measured by \citet{lockwood14}. } } \end{table} The dwarf planet Haumea \citep{brown05} is an extremely fast rotator with density higher than other objects in the Kuiper belt; it is consistent with a body dominated by rock \citep{rabinowitz06,lacerda08,lellouch10,kondratyev16}. Visible and infrared light curve fits \citep{lockwood14} find the body consistent with a rapidly rotating oblong Jacobi ellipsoid shape in hydrostatic equilibrium with axis ratios listed in Table \ref{tab:haumea} (also see \citealt{lellouch10}) and a density of $\rho = 2.6$~g~cm$^{-3}$. For discussion on formation scenarios for the satellite system see \citet{leinhardt10,schlichting09,cuk13}. Parameters based on Haumea and Hi'iaka are listed in Table \ref{tab:haumea}. The pressures in the body at depth for a body as massive as Haumea would lead to ductile flow giving long-term deformation allowing the body to approach a figure of equilibrium (a Jacobi ellipsoid). Even if Haumea's shape is consistent with a hydrostatic equilibrium figure, on short timescales the body should behave elastically. For tidal evolution, the relevant tidal frequency is $\omega \sim 2 \sigma_H \sim 10^{-3}$Hz, comparable to vibrational normal mode frequencies in the Earth. Our simulations do not allow ductile flow on long timescales, but can approximate the faster tidal deformations if we model the body as a stiff elastic body with its current shape. We ran a simulation in the LR random lattice series with axis ratios $b/a=0.8$ and $c/a=0.5$, consistent with measurements for Haumea. The LR series of simulations has more particles than the R series and is discussed in more detail in section \ref{ap:num}. From the simulation we measure $\dot a_o/a_s = 2.04$ or drift rate approximately twice that of the equivalent volume sphere. Equation \ref{eqn:tri_line} predicts a value 2.034, consistent with the numerical measurement, whereas equation \ref{eqn:scale_tri} gives 2.39. In their section 4.3.1, \citet{ragozzine09} speculated that using the volumetric radius leads to an underestimate of the tidal evolution. However for the axis ratio of Haumea $b/a \approx 0.8$ and $c/a \approx 0.52$, we find here that the drift rate would only be about twice as fast as estimated using the volumetric radius. \begin{figure} \includegraphics[width=3.3in]{haumea_ice_comparison.png} \caption{We compare the drift rate in orbital semi-major axis for two simulations in the LR series and with axis ratios $b/a=0.8$ and $c/a=0.6$, similar to Haumea. The blue points show a simulation of a homogeneous body, whereas the black points show a simulation with weaker ends (where springs have a lower spring constant), mimicking a rocky body with icy ends. The lines show linear fits measuring the secular drift rate. The $x$ axis shows time in units of $t_g$ (equation \ref{eqn:tgrav}) and the $y$ axis shows semi-major axis in units of volumetric radius, $R_v$, measured from the initial value. The drift rate of the body with soft ends is about 5.2 times faster than the homogeneous ellipsoid with the same axis ratios and about 10 times faster than the equivalent volume homogeneous sphere. \label{fig:icy}} \end{figure} \citet{kondratyev16} proposed that stresses between icy shell and core and associated relaxation would cause ice to accumulate at the ends of Haumea. He proposed that the icy ends could separate forming the two icy satellites Namaka and Hi'iaka. Estimates for the fraction of ice in a differentiated Haumea range from 7\% \citep{kondratyev16} to 30\% \citep{probst15}. Young's modulus of ice is estimated to be a few GPa (\citealt{nimmo06}; see \citealt{collins10} for a review) and this is about 10 times lower than the Young's modulus for rocky materials. To explore the affect of softer icy ends on the semi-major axis drift rate, we ran a simulation of a body that is not homogeneous. Using the same axis ratios of $b/a=0.8$ and $c/a=0.5$ and parameters of the LR series, we reduced the spring constants to 1/10th the value in the body core at radii greater than 1 (in units of volumetric radius) from the body center. About 20\% of the springs have reduced spring constants. This has the effect of lowering the simulated Young's modulus at the ends of the ellipsoid by a factor of 10. We did not vary the density as the difference in density between ice and rock is much lower than their difference in elastic modulus. The spring damping parameter $\gamma$ does not vary, so $\tau$, the viscoelastic relaxation time-scale, and tidal frequency, $\bar \chi$, are the same in both regions. The measurements of semi-major axis for this simulation and for the homogeneous one with the same axis ratios are shown in Figure \ref{fig:icy} along with linear fits that measure the secular drift rate. We measured the drift rate in semi-major axis in this simulation, finding that it is about 5.2 times faster than the homogeneous ellipsoid with the same axis ratios and 10 times faster than the equivalent homogeneous sphere. Even a small fraction of softer material can significantly affect the simulated drift rates. Perhaps this should have been expected based on the strong sensitivity to elastic anisotropy that we inferred from the cubic lattice model simulations. The classical tidal formula for the tidal drift rate in semi-major axis \begin{equation} \frac{\dot a_o}{n a_o} = \frac{3 k_{2H}}{Q_H} \frac{M_*}{M_H} \left( \frac{R_{vH}}{a_o} \right)^5 \end{equation} (e.g., \citealt{M+D}) for perturbing object $M_*$ (here Hi'iaka) due to tidal dissipation in the spinning body $M_H$, where $k_{H}$ and $Q_H$ are the Love number and dissipation factor for Haumea. We can describe corrections to the tidal drift rate by multiplying the right hand side by a parameter $f_{corr}>1$. The above equation implies that $\dot a_o \propto a_o^{-5.5}$ (taking into account dependence on $n a_o \propto a_o^{-1/2}$). We integrate the above equation for Haumea to estimate the time it takes Hi'iaka to tidally drift outwards to its current semi-major axis. Putting unknowns on the left hand side \begin{equation} \frac{k_{2H}}{Q_H} f_{corr} \sim \frac{1}{n \tau_a} \left( \frac{a_{Hi}}{R_v} \right)^5 \frac{2}{39} q^{-1} \sim 0.1 \end{equation} with mass axis ratio $q =M_{Hi}/M_{H}$ and where we have used values from Table \ref{tab:haumea} and an age $\tau_a = 4$ Gyr for the timescale over which tidal migration is taking place. Here $n$ and $a_{Hi}$ are the mean motion and orbital semi-major axis of Hi'iaka at its current location. If $k_{2H}$ for Haumea is as large as 0.01 (at the border of what would be consistent with rigidity for rocky material) and we use corrections for shape and composition increasing the drift rate by $f_{corr} = 10$ then we have equality only if the dissipation parameter is large; $Q \sim 1$. We conclude that it is unlikely that Hi'iaka alone tidally drifted to its current location even if the tidal drift rate is larger by a factor of 10 than estimated using the equivalent volume rocky sphere. One explanation for the origin of Haumea's satellites and compositional family is a collisional disruption of a past large moon of Haumea \citep{schlichting09}. The `ur-satellite' would have formed closer to Haumea and because of its large mass, could have migrated more quickly than Hi'iaka outward during the lifetime of the Solar system. The failure of our enhanced tidal drift rate estimate to account for Hi'iaka's current position would suport the `ur-satellite' proposal (also see discussion by \citealt{cuk13}).

1607.08591

1607

1607.06773_arXiv.txt

We present direct estimates of the mean sky brightness temperature in observing bands around 99GHz and 242GHz due to line emission from distant galaxies. These values are calculated from the summed line emission observed in a blind, deep survey for spectral line emission from high redshift galaxies using ALMA (the 'ASPECS' survey). In the 99 GHz band, the mean brightness will be dominated by rotational transitions of CO from intermediate and high redshift galaxies. In the 242GHz band, the emission could be a combination of higher order CO lines, and possibly [CII] 158$\mu$m line emission from very high redshift galaxies ($z \sim 6$ to 7). The mean line surface brightness is a quantity that is relevant to measurements of spectral distortions of the cosmic microwave background, and as a potential tool for studying large-scale structures in the early Universe using intensity mapping. While the cosmic volume and the number of detections are admittedly small, this pilot survey provides a direct measure of the mean line surface brightness, independent of conversion factors, excitation, or other galaxy formation model assumptions. The mean surface brightness in the 99GHZ band is: $T_B = 0.94\pm 0.09\mu$K. In the 242GHz band, the mean brightness is: $T_B = 0.55\pm 0.033\mu$K. These should be interpreted as lower limits on the average sky signal, since we only include lines detected individually in the blind survey, while in a low resolution intensity mapping experiment, there will also be the summed contribution from lower luminosity galaxies that cannot be detected individually in the current blind survey.

Intensity mapping of the cumulative CO and other millimeter and submillimeter line emission from early galaxies has been proposed as a new means to probe very large-scale structures in the distant Universe (Carilli 2011; Gong et al. 2011; Gong et al. 2012; Yue et al. 2015). Intensity mapping entails low spatial and spectral resolution imaging of the sky to obtain the mean brightness due to the cumulative emission from myriad discrete cosmic sources. While interferometric arrays like ALMA, the JVLA, and NOEMA, can detect CO and [CII] 158$\mu$m (and in cases of high luminosity sources, other lines), from individual galaxies at high redshift, the fields of view are very small, and the integration times are long. These telescopes are inadequate for measuring the galaxy distribution on the very large scales relevant to cosmological questions, such as the Baryon Acoustic Oscillations at intermediate redshifts, or the large-scale distribution of galaxies that reionize the Universe. The latter is of particular interest for cross correlation studies with very wide field, low-resolution HI 21cm images of the intergalactic medium during cosmic reionization (Lidz et al. 2011). The integrated millimeter and submillimieter line emission from early galaxies has also been recognized as a possible significant contaminant of measurements of the spectral and spatial fluctuations of the cosmic microwave background (CMB; Righi et al. 2008a,b; Chluba \& Sunyaev, 2012; de Zotti et al. 2015; Mashian et al. 2016). For example, modeling suggests (Mashian et al. 2016) that the integrated CO line emission could be significantly higher than the primordial spectral distortions due to other cosmological effects (e.g., Chluba \& Sunyaev, 2012; Sunyaev \& Khatri, 2013; Tashiro 2014), and may be measurable with next generation instruments like the Primordial Inflation Explorer (PIXIE; Kogut et al. 2014).\footnote{PIXIE is a space observatory concept to map the CMB over the frequency range 30GHz to 6THz, one goal of which is to constrain the average CMB energy spectrum with much greater accuracy than FIRAS.} Numerous calculations have been done to predict the mean sky brightness due to emission lines from CO at intermediate and high redshift, and [CII] 158$\mu$m emission at very high redshift (see section 2). These predictions are based on either empirical estimates using proxies for the line emission, such as the cosmic star formation rate density, or large scale cosmological simulations of galaxy formation, with recipes to relate proxy measurements or simulated properties to line luminosities. In this paper, we present direct measurements of the summed line luminosity from individual sources in bands around 99 GHz and 242 GHz. These measurements are based on the ASPECS program, corresponding to a broad band spectral line deep field of the UDF at 1.25mm and 3mm (Decarli et al. 2016; Aravena et al. 2016a; Walter et al. 2016). From these measurement, we derive the mean brightness temperature at a given observing frequency due to high redshift galaxies. As a pilot study with ALMA, the fields are necessarily small, and the number of galaxies few. However, the measured quantity is direct: line emission from early galaxies. Hence, no modeling or conversion factors are required.

\subsection{Limits} As a pilot ALMA study, we reemphasize that the volumes in question are small, as are the number of detections. Hence, our conclusions and uncertainties are dominated by cosmic variance and simple shot noise (Poisson statistics). Aravena et al. (2016b) consider the issue of cosmic variance in the context of our particular field. Based on the drop-out galaxy counts, and the bright submm source counts, this bias might be as large as a factor two (low). On the other hand, consideration of the contribution of faint submm continuum sources to the cosmic infrared background, based on our deeper ASPECS ALMA data, suggests a factor closer to unity (Aravena et al. 2016b). Regardless, since this is a direct survey of the observable in question, namely, mean brightness due to line emission from distant galaxies at a given observing frequency, the results remain of interest in general progress toward millimeter line intensity mapping, and a factor two uncertainty is inconsequential for our analysis in section 4.2. Our measurements are also lower limits, since we only sum lines detected. We do not extrapolate to, e.g., lower or higher luminosity galaxies using an assumed luminosity function. Considering CO (the dominant contributor at 99GHz, certainly, and likely at 242 as well), our detection threshold was set in order to reach what may be the 'knee' in the CO luminosity function at the primary redshifts to which our survey is most sensitive ($z \sim 1$ to 3). This estimation was based on both numerical simulations and extrapolations of CO emission properties of high redshift galaxies from, e.g., measures of dust luminosities or star formation rates (see Decarli et al. 2016a for more details). If the CO luminosity function is relatively flat at low luminosities, and steep at high luminosities, then galaxies around the knee of the curve dominate the overall luminosity. For example, using the Popping et al. (2016) and Lagos et al. (2012) CO luminosity functions and our limits at 99GHz, we estimate that we should be detecting between 40\% and 70\% of the total CO luminosity (Decarli et al. 2016) in this dominant redshift range. \subsection{Comparison to predictions and CMB spectral distortions} As stated in Section 1, millimeter line intensity mapping experiments will have broad impact, from studies of galaxy formation to the Baryon Acoustic Oscillations. In this section we consider in some detail our results in the context of one topical area that has seen considerable attention recently, namely, the spectral distortions of the CMB. In section 2, we reviewed the predictions for the line brightness at 99GHz and 242 GHZ based on phenomenological calculations using on proxies for the line luminosity (such as the cosmic star formation density), or numerical simulations of galaxy formation. Predictions vary significantly, but range from $\sim 1\mu$K to $10\mu$K, in the frequency ranges being considered. To within the uncertainties inherent in small volume surveys, our direct measurements of $T_B = 0.94\pm 0.09\mu$K at 99GHz and $T_B = 0.55\pm 0.033\mu$K at 242GHz, argue for the faint end of these predictions, although we again emphasize that these should be treated as lower limits. How do our measurements then compare to, for instance, the expected distortions in the CMB spectrum due to early energy release, and to the expected sensitivity of planned CMB spectral distortion experiments? As a benchmark for experimental sensitivity, we adopt the current parameters being considered for PIXIE (Kogut et al. 2014; 2011), using the 15 GHz spectral resolution for the proposed experiment. Considering the expected sky brightness contributions, we focus on the more cosmologically relevant predictions, relating to recombination and reionization. We note that there are other potentially significant foregrounds, in particular, Galactic and extragalactic thermal emission from warm dust, and synchrotron emission. Kogut et al. (2014) review the relative magnitude of these contributions. The thermal emission from warm dust, in particular, is calculated to be an order of magnitude, or more, stronger than the summed millimeter line emission considered herein. However, the spectral behavior of the dust emission is considered to be well understood, and should be well modelled, and removed, using spectral fitting algorithms over a broad frequency range. Herein, we focus on the millimeter and sub-millimeter line emission, given that this is our measured quantity, and compare it to the predicted cosmological signals. Additional discussion of foregrounds can be found in, e.g., de Zotti et al. (2015). \begin{figure} \centering \includegraphics[width=\columnwidth]{./DT_spec.pdf} \caption{Comparison of different CMB distortion signals (negative branches of the signals are marked) with the millimeter line limits reported in this paper. The low-redshift distortion created by reionization and structure formation is close to a $y$-distortion with $y\simeq 2\times 10^{-6}$. Contributions from the hot gas in low mass halos give rise to a noticeable relativistic temperature correction. For the damping signal, we plot a $\mu$-distortion with $\mu=2\times 10^{-8}$. The cosmological recombination radiation was computed using {\tt CosmoSpec}. The estimated effective sensitivity ($\Delta I_\nu \approx 5\, {\rm Jy~ sr^{-1}}$) of PIXIE is shown for comparison (dotted line).} \label{fig:one} \end{figure} In Fig.~3, we show a comparison of various distortion signals, along with the line limits derived herein. We focus on {\it guaranteed} distortions within $\Lambda$CDM (see Chluba, 2016, for most recent overview), some of which should be detectable with future experiments, at least in terms of raw sensitivity (Kogut et al. 2011). A wider range of range of energy release processes (e.g., decaying particle scenarios) is discussed in Chluba 2013 and Chluba \& Jeong, 2014. The largest CMB expected spectral distortion is created at low-redshift by the reionization and structure formation process (Sunyaev \& Zeldovich, 1972; Hu et al., 1994a). This signal is close to a pure Compton-$y$ distortion (Zeldovich \& Sunyaev, 1969) caused through partial up-scattering of CMB blackbody photons by hot electrons yielding a $y$-parameter $y\simeq 2\times 10^{-6}$ (e.g., Refregier et al. 2000; Hill et al., 2015). Contributions from the hot gas ($\simeq 1{\rm keV}$) residing in low mass halos also give rise to a noticeable relativistic temperature correction, which could be used to constrain the average temperature of baryons at low redshifts (Hill et al., 2015). While the relativistic correction signal requires a removal of the integrated CO emission, the non-relativistic $y$-distortion contribution should be less affected and already separable using multi-frequency capabilities of future experiments. Another inevitable distortion is created by the dissipation of small-scale fluctuations in the primordial photon-baryon plasma (Sunyaev \& Zeldovich, 1970; Daly, 1991; Hu et al., 1994b; Chluba, Khatri, \& Sunyaev 2012) due to Silk damping. We illustrate the $\mu$-distortion (Sunyaev \& Zeldovich, 1970) contribution of this signal using $\mu=2\times 10^{-8}$, which is close to the value expected for the $\Lambda$CDM cosmology (Chluba, 2016). A $\mu$-distortion can only be created in dense and hot environments present in the early Universe at $z\gtrsim 5\times 10^4$ (Burigana et al., 1991; Hu \& Silk, 1993). By detecting this signal one can probe the amplitude of perturbations at scales far smaller than those seen in the CMB anisotropies, delivering another independent way to test different inflation models (e.g., Chluba, Khatri, \& Sunyaev 2012; Chluba, Erickcek \& Ben-Dayan, 2012; Dent et al., 2012; Clesse et al., 2014). Finally, we show the cosmological hydrogen and helium recombination radiation emitted at $z\simeq 10^3$ (Zeldovich et al., 1968; Peebles, 1968; Dubrovich, 1975; Kholupenko et al., 2005; Rubi{\~n}o-Mart{\'{\i}}n et al., 2006; Chluba \& Sunyaev 2006), which was computed using {\tt CosmoSpec} (Chluba \& Ali-Ha{\"i}moud, 2016). This signal could provide an independent way to constrain cosmological parameters and directly map the recombination history (Sunyaev \& Chluba, 2009). It is unpolarized and its unique spectral variability is very hard to mimic by other foregrounds or instrumental effects.\footnote{The expected distortions due to annihilating dark matter (McDonald et al., 2001; Chluba 2010; Chluba \& Sunyaev, 2012; Chluba 2013) and the differences in the cooling of baryons relative to CMB photons during cosmic expansion (Chluba, 2005; Chluba \& Sunyaev, 2012) were not illustrated here.} The latter two effects cause fractional spectral distortions in the range of 10$^{-9}$ to 10$^{-8}$, implying observed brightness temperature perturbations $\Delta T_{\rm B}\simeq 3{\rm nK}-30{\rm nK}$, well below the contribution of the mean line brightness measured herein. Thus, beyond doubt, an extraction of these primordial distortions will be very challenging, requiring sophisticated foreground removal techniques, unprecedented control of systematics, broad spectral coverage and high sensitivity multi-frequency capabilities. To successfully remove the integrated millimeter and submillimeter line emission, it will be advantageous to exploit the synergies between future CMB distortion measurements and observations similar to those presented here. Given the importance of the primordial distortion signals to studies of early-universe physics, this direction is highly relevant. As ALMA attains full capability, spectral deep fields will become more efficient and effective, eventually encompassing areas of tens of square arcminutes. Our pilot studies have already shown the impact of such measurements over a broad range of problems in modern astrophysics and cosmology. In parallel, the Jansky Very Large Array is exploring similar deep spectral searches at 30GHz (eg. Lentati et al. 2015; Riechers et al. in prep.), while the advent of high frequency spectral cameras on the Green Bank Telescope provide a sensitive platform for wide field spectral searches (Sieth et al. 2016). In the long term, a 'Next Generation Very Large Array,' operating between 20GHz and 115GHz with octave, or broader, bandwidth receivers and ten times the collecting area of ALMA and the JVLA, has the potential to revolutionize blind searches for molecular gas in the early Universe (Carilli et al. 2015; Casey et al. 2015).

1607.06773

1607

1607.01365_arXiv.txt

The extension of the singular perturbative approach to the second order is presented in this paper. The general expansion to the second order is derived. The second order expansion is considered as a small correction to the first order expansion. Using this approach it is demonstrated that the second order expansion is reducible to a first order expansion via a re-definition of the first order pertubative fields. Even if in practice the second order correction is small the reducibility of the second order expansion to the first order expansion indicates a degeneracy problem. In general this degeneracy is hard to break. A useful and simple second order approximation is the thin source approximation which offers a direct estimation of the correction. The practical application of the corrections derived in this paper are illustrated by using an elliptical NFW lens model. The second order pertubative expansion provides a noticeable improvement, even for the simplest case of thin source approximation. To conclude it is clear that for accurate modelisation of gravitational lenses using the perturbative method the second order perturbative expansion should be considered. In particular an evaluation of the degeneracy due to the second order term should be performed, for which the thin source approximation is particularly useful.

The singular perturbative method is a non parametric approach to gravitational lenses offering a direct relation between the description of the lens and the observations. The direct relation between the lens and the data minimize the degeneracy problems generally encountered in gravitational lenses modeling (see for instance ~\cite{Saha2006}, ~\cite{Wucknitz2002}, ~\cite{Chiba2002}). In a series of papers ~\cite{Alard2007} ~\cite{Alard2009} ~\cite{Alard2010} the first order singular perturbative method was considered. Let's first recall the basics of the first order method. We consider a perturbation of the perfect ring situation. A perfect ring is obtained when a point source is at the center of a circular potential. The images of the central point source is an infinity of points situated on a circle. The radius of this circle is the Einstein radius associated with the circular potential. For simplicity the Einstein radius is reduced to unity by adopting a proper set of distance units. The introduction of a non circular perturbation to the circular potential results in the breaking of the circle with the consequence that the central point has now a finite number of images in the vicinity of the circle. In practice the source itself is not reduced to a point but has a finite size which is of the order of the potential perturbation. Additionally the source may not be exactly at the center of the circular potential and as a consequence has an impact parameter which is also of the order of the potential perturbation which we call $\epsilon$, with $\epsilon \ll 1$. Using polar coordinates ($r$,$\theta$) in the lens plane, the potential reads: \begin{equation} \phi(r,\theta)=\phi_0(r)+\epsilon \psi(r,\theta) \label{pot_def} \end{equation} The lens equation relating the lens plane coordinates ${ \bf r}$ to the source plane coordinates ${\bf r_S}$ reads: \begin{equation} {\bf r_S} = {\bf r} -\nabla \phi \label{lens_eq} \end{equation} The radial deviation from the circle is of the same order as the potential perturbation, thus $r=1+\epsilon dr$. By inserting Eq. ~\ref{pot_def} in the lens equation and developing to the first order in $\epsilon$ we obtain a set of equations already presented in \cite{Alard2007} \begin{equation} {\bf r_S} = \left(\kappa_2 \ dr-f_1 \right) {\bf u_r} - \frac{d f_0}{d \theta }\bf {u_{\theta}} \label{pert_1} \end{equation} And: \begin{equation} f_1=\left[\frac{\partial \psi}{\partial r} \right]_{r=1} \ \ f_0=\psi(1,\theta) \ \ \kappa_2=1-\frac{d^2 \phi_0}{d r^2} \end{equation} Considering that the source has an impact parameter, ${\bf r_0}=(x_0,y_0)$ it is useful to re-write Eq. ~\ref{pert_1} using the variable ${\bf r_S={\tilde r_S}+r_0}$: \begin{equation} {\bf \tilde r_S} = \left(\kappa_2 \ dr-\tilde f_1 \right) {\bf u_r} - \frac{d \tilde f_0}{d \theta }\bf {u_{\theta}} \label{pert_1i} \end{equation} With: \begin{equation} \tilde f_i=f_i+x_0 \cos(\theta) + y_0 \sin(\theta) \ \ , \ i=0,1 \end{equation}

It is relatively simple to estimate the second order perturbative expansion as a correction of the first order expansion. In particular the correction in the thin source limit is simple and provide a noticeable improvement over the first order pertubative expansion. The full iterative second order correction converge to the second order perturbative correction but in most cases provides only a small additional improvement with respect to the thin source approximation. Additionally it is interesting to note that larger sources can always be de-composed in a number of thinner sources for which the thin source approximation is valid. Another important issue is the problem of the degeneracy of the second order correction. Even if in most case the correction is small the problem of the possible degeneracy of first order expansion should be addressed. For an evaluation of the amplitude of the degenerate term the thin source approximation should be particularly useful as it offers a direct estimation. In some particular application when the degeneracy of the second order term can be broken the full estimation of the second order expansion should be useful.

1607.01365

1607

1607.07296_arXiv.txt

We investigate the matter density perturbation $\delta_m$ and power spectrum $P(k)$ in the running vacuum model (RVM) with the cosmological constant being a function of the Hubble parameter, given by $\Lambda = \Lambda_0 + 6 \sigma H H_0+ 3\nu H^2$, in which the linear and quadratic terms of $H$ would originate from the QCD vacuum condensation and cosmological renormalization group, respectively. Taking the dark energy perturbation into consideration, we derive the evolution equation for $\delta_m$ and find a specific scale $d_{cr}=2 \pi/k_{cr}$, which divides the evolution of the universe into the sub and super-interaction regimes, corresponding to $k \ll k_{cr}$ and $k \gg k_{cr}$, respectively. For the former, the evolution of $\delta_m$ has the same behavior as that in the $\Lambda$CDM model, while for the latter, the growth of $\delta_m$ is frozen (greatly enhanced) when $\nu + \sigma >(<)0$ due to the couplings between radiation, matter and dark energy. It is clear that the observational data rule out the cases with $\nu<0$ and $\nu + \sigma <0$, while the allowed window for the model parameters is extremely narrow with $\nu, |\sigma| \lesssim \mathcal{O}(10^{-7})$.

\label{sec:introduction} It is well-known that the Type-Ia supernova observations (Riess et al. \citep{Riess:1998cb}; Perlmutter et al. \citep{Perlmutter:1998np}) have revealed the late-time accelerating expansion of our universe. To realize the accelerating universe, it is necessary to introduce a negative pressure fluid to the gravitational theory, referred to as ``Dark Energy'' (Copeland et al. \citep{Copeland:2006wr}), while the simplest scenario is to have the cosmological constant $\Lambda$, such as the $\Lambda$CDM model. Currently, the $\Lambda$CDM model perfectly fits the observational data, but leaves several difficulties, such as the ``fine-tuning" (Weinberg \citep{Weinberg:1988cp}; Weinberg \citep{WBook}) and ``coincidence'' (Ostriker and Steinhardt \citep{Ostriker:1995rn}; Arkani-Hamed et al. \citep{ArkaniHamed:2000tc}) problems. In this work, we are interested in the running vacuum model (RVM), which has been used to solve the ``coincidence'' problem (Ozer and Taha~\citep{Ozer:1985ws}; Carvalho et al. \citep{Carvalho:1991ut}; Lima and Maia \citep{Lima:1994gi}; Lima and Trodden \citep{Lima:1995ea}; Overduin and Cooperstock \citep{Overduin:1998zv}; Dymnikova and Khlopov \citep{Dymnikova:2001ga}; Carneiro and Lima \citep{Carneiro:2004iz}; Bauer \citep{Bauer:2005rpa}; Shapiro et al. \citep{Shapiro:2004ch}; Alcaniz and Lima\citep{Alcaniz:2005dg}; Barrow and Clifton \citep{Barrow:2006hia}; Shapiro and Sola \citep{Shapiro:2009dh}; Geng and Lee \citep{Geng:2016epb}; Geng et al. \citep{Geng:2016dqe}). In this model, the cosmological constant evolves in time and decays to radiation and matter in the evolution of the universe, leading to the same order of magnitude for the energy densities of dark energy and dark matter. Its observational applications have been also extensively explored in the literature (Espana-Bonet et al. \citep{EspanaBonet:2003vk}; Tamayo et al. \citep{Tamayo:2015qla}). Additionally, it has been shown that the RVM can fit various observational data, indicating that this scenario is good in describing the evolution history of our universe (Sola \citep{Sola:2016vis}; Sola et al. \citep{Sola:2015wwa}; Sola et al. \citep{Sola:2016jky}; Sola et al. \citep{Sola:2016ecz}). In our study, we will concentrate on the specific model with $\Lambda= \sum\limits_{i=0}^{2} \lambda_i H^i$ (Borges and Carneiro \citep{Borges:2005qs}; Borges et al. \citep{Borges:2007bh}; Carneiro et al. \citep{Carneiro:2007bf}; Zimdahl et al. \citep{Zimdahl:2011ae}; Sola \citep{Sola:2013gha}; Sola and Gomez-Valent \citep{Sola:2015rra}), in which the quadratic term, $\lambda_2 H^2$, might come from the quantum effects induced by the cosmological renormalization group (Alcaniz et al. \citep{Alcaniz:2012mh}; Costa et al. \citep{Costa:2012xw}; Sola \citep{Sola:2014tta}; Gomez-Valent et al. \citep{Gomez-Valent:2014rxa}), while the linear term, $\lambda_1 H$, would originate from the theory with the QCD vacuum condensation associated with the chiral phase transition (Schutzhold \citep{Schutzhold:2002pr}; Banerjee et al.\citep{Banerjee:2003fg}; Klinkhamer and Volovik \citep{Klinkhamer:2009nn}; Ohta \citep{Ohta:2010in}; Cai et al. \citep{Cai:2010uf}). When it comes to the decaying dark energy model, it is reasonable to consider not only the background evolution equations but also the density perturbation of dark energy. We follow the same method in the references (Fabris et al. \citep{Fabris:2006gt}; Borges et al. \citep{Borges:2008ii}) to rewrite dark energy as a function of a Lorentz scalar $\nabla_{\mu} U^{\mu}$, where $U^{\mu}=dx^{\mu}/\sqrt{-ds^2}$ is the four-velocity. Based on such an expression, we examine the matter density perturbation $\delta_m$ and power spectrum $P(k)$ in the linear perturbation theory of gravity. Note that in the literature (Fabris et al. \citep{Fabris:2006gt}), the matter density perturbation evolves from $z=1100$ (the recombination era) to $z=0$ (the present), where the initial conditions are taken from the $\Lambda$CDM limit with the BBKS transfer function. However, the density perturbation of the RVM may influence the evolution of the matter density perturbation in the high redshift regime. We take the scale invariance initial conditions at the very early time of the universe, in which all the perturbation modes are at the super-horizon scale with the same behavior as that in the $\Lambda$CDM model. Then, we analyze the properties in the sub and super-interaction scales with the allowed ranges for the model parameters discussed. This paper is organized as follows: We briefly introduce the running vacuum model in Sec.~\ref{sec:model}. We derive the linear perturbation equations with the synchronous gauge and the evolution property of the matter density perturbation in Sec.~\ref{sec:density-pert}. In Sec.~\ref{sec:mpk}, we show the evolutions of $\delta_m$ and $P(k)$. Our conclusions are presented in Sec.~\ref{sec:conclusion}.

\label{sec:conclusion} We have studied the matter density perturbation $\delta_m$ and matter power spectrum $P(k)$ in the RVM with $\Lambda = \Lambda_0 + 6 \sigma H H_0+ 3\nu H^2$. By rewriting $\Lambda$ to be a function of the covariant derivative of four-velocity as $\Lambda = \Lambda_0+ 2 \sigma \nabla_{\mu} U^{\mu} + \nu (\nabla_{\mu} U^{\mu})^2/3$, we explicitly derive the linear perturbation equations for matter and radiation. The dark energy perturbation $\delta_{\Lambda}$ can be expressed by $\theta$ and $h$, indicating that $\delta_{\Lambda}$ directly couples to $\delta_m$ and $\theta_m$. We have shown that the growth of $\delta_m$ can be separated into the sub and super-interaction regimes of $\vert \tilde{k}^2 \vert /a^2 \ll H^2$ and $\vert \tilde{k}^2 \vert /a^2 \gg H^2$, respectively. In the former regime, the interactions between dark energy and matter are sub-dominated to the evolutions of $\delta_m$ and $\theta_m$, and the growth of $\delta_m$ behaves the same as that in the $\Lambda$CDM model. In the later one, the decaying $\Lambda$ drags the evolution of $\delta_m$, and $P(k)$ is suspended (sharply increased) when $\nu+\sigma H_0 / H >(<)\,0$. The RVM with $\nu<0$ or $\nu+\sigma <0$ is clearly ruled out by the divergences of the physical quantities, $\delta_m$ and $\theta_m$. We have also found that the model parameters are strongly constrained to be $\nu>0$ and $\nu+\sigma >0$ with $\nu$ and $|\sigma| \lesssim \mathcal{O}(10^{-7})$. The perturbed RVM modifies not only the growth of $\delta_m$ but the evolution of $\theta_m$. In the $\Lambda$CDM model, the cold dark matter rests on the comoving frame, i.e., $\theta_m \rightarrow 0$, but the behavior of $\theta_m$ in the RVM scenario is totally different. In the super-interacting regime, $\delta_m$ is frozen, but $\theta_m$ is enhanced to be a non-zero value, indicating that the massive cold dark matter is heated up by the decaying dark energy. This kind of the enhancement of $\theta_m$ might significantly increase the velocity of dark matter. To realize this effect, we have to further investigate physics at the scale of the dark matter halo, at which the linear perturbation theory is no longer valid, and the non-perturbative calculation is needed.

1607.07296

1607

1607.02150_arXiv.txt

We use simple toy models of far-IR dust emission to estimate the accuracy to which the polarization of the cosmic microwave background can be recovered using multi-frequency fits, if the parametric form chosen for the fitted dust model differs from the actual dust emission. Commonly used approximations to the far-IR dust spectrum yield CMB residuals comparable to or larger than the sensitivities expected for the next generation of CMB missions, despite fitting the combined CMB $+$ foreground emission to precision 0.1\% or better. The Rayleigh-Jeans approximation to the dust spectrum biases the fitted dust spectral index by $\Delta \beta_d = 0.2$ and the inflationary B-mode amplitude by $\Delta r = 0.03$. Fitting the dust to a modified blackbody at a single temperature biases the best-fit CMB by $\Delta r > 0.003$ if the true dust spectrum contains multiple temperature components. A 13-parameter model fitting two temperature components reduces this bias by an order of magnitude if the true dust spectrum is in fact a simple superposition of emission at different temperatures, but fails at the level $\Delta r = 0.006$ for dust whose spectral index varies with frequency. Restricting the observing frequencies to a narrow region near the foreground minimum reduces these biases for some dust spectra but can increase the bias for others. Data at THz frequencies surrounding the peak of the dust emission can mitigate these biases while providing a direct determination of the dust temperature profile.

Polarization of the cosmic microwave background (CMB) provides a critical test for models of inflation. The primary signature is is a parity-odd curl component in the polarization on angular scales of a few degrees or larger \citep{kamionkowski/etal:1997, seljak/zaldarriaga:1997}. The amplitude of this ``B-mode'' signal depends on the inflationary potential \begin{equation} V^{1/4} = 1.06 \times 10^{16} ~{\rm GeV} \left( \frac{r}{0.01} \right)^{1/4} \label{potential_eq} \end{equation} where $r$ is the power ratio of the tensor (gravitational) perturbations to scalar (density) fluctuations \citep{turner/white:1996}. If inflation results from Grand Unified Theory physics (energy $\sim 10^{16}$ GeV), the B-mode amplitude should be in the range 1 to 100 nK. Signals at this amplitude could be detected by a dedicated polarimeter, providing a critical test of a central component of modern cosmology. Detecting the inflationary signal will be challenging. A primary concern is confusion from astrophysical foregrounds. We view the CMB through a screen of diffuse Galactic emission originating within different components of the interstellar medium. Figure \ref{foreground_fig} compares the inflationary B-mode signal to polarized Galactic foregrounds. Synchrotron emission from relativistic cosmic ray electrons accelerated in the Galactic magnetic field dominates the diffuse radio continuum at low frequencies. For a power-law distribution of cosmic ray energy $N(E) \sim E^{-p}$, the synchrotron intensity is also a power law, \begin{equation} I_s(\nu) = A_s \left(\frac{\nu}{\nu_s} \right)^{\beta_s}, \label{synch_power_law} \end{equation} \noindent where $\nu$ is the observing frequency, $\beta_s = (1-p)/2$ is the spectral index, and $A_s$ is the amplitude defined relative to reference frequency~$\nu_s$ \citep{rybicki/lightman:1979}. The measured values $2.6 < p < 3.2$ for the cosmic ray energy spectrum correspond to synchrotron spectral index $-1.1 < \beta_s < -0.8$, in reasonable agreement with radio data \citep{strong/etal:2007, jaffe/etal:2011, kogut:2012, bennett/etal:2013, planck:2015:10}. The synchrotron spectrum may contain additional features (curvature, spectral break), but evidence for such features is restricted to frequencies below 20 GHz and is not considered here. Dust is the dominant foreground at high frequencies. Dust grains in the interstellar medium absorb optical and UV photons and re-radiate the energy in the far-infrared. The resulting spectrum is often empirically modeled as a sum of modified blackbodies with power-law emissivities, \begin{equation} I_d(\nu) = \sum_i \epsilon_i B_\nu(T_i) \left( \frac{\nu}{\nu_d} \right)^{\beta_{d\,i}} \label{dust_greybody_eq} \end{equation} where $B_\nu(T)$ is the Planck intensity at temperature $T$, and the emissivity $\epsilon$ and spectral index $\beta_d$ are defined at reference frequency $\nu_d$. Fitting the high-latitude dust cirrus to a single modified blackbody returns values $T_d~=~20$ K and $\beta_d = 1.6$ \citep{planck:2015:22}. The B-mode signal is fainter than the Galactic foregrounds at all frequencies. Current measurements limit primordial B-modes to amplitude $r < 0.07$ at 95\% confidence \citep{bicep/planck:2015}. Distinguishing primordial B-modes from Galactic foregrounds at this level requires subtracting the foregrounds to few percent accuracy or better. The next generation of CMB polarimeters anticipates sensitivities $r < 0.001$. Measurements at this level require foreground subtraction with sub-percent accuracy. Despite their importance, diffuse Galactic foregrounds at millimeter wavelengths are poorly constrained. The observed dust emission depends on a number of factors including the grain size distribution, chemical and physical composition of the dust grains, competing emission mechanisms within a grain, and the three-dimensional distribution of the dust population irradiated by the stellar UV/optical field. None of these are known in detail. Lacking a detailed physical model, CMB analyses typically use a purely phenomenological model for the dust, treating it as the sum of one or two modified blackbody components along each line of sight. In this paper, we use simple extensions to commonly used models of far-IR dust emission to estimate the systematic error (bias) in the B-mode amplitude $r$ due to the use of such phenomenological models. \begin{figure}[t] \vspace{-2mm} % \includegraphics[width=2.6in, angle=90]{pol_foregrounds_apj_2015.eps} \caption{Frequency spectra of the CMB and polarized foregrounds. The grey band shows 0.01$<$$r$$<$0.1 for the primordial inflation signal. Colored bands show the synchrotron and dust foregrounds for the cleanest 50\% and 75\% of the sky. The inflationary signal is fainter than Galactic foregrounds at all frequencies, requiring accurate models for foreground subtraction. } \label{foreground_fig} \end{figure}

Parametric modeling of the diffuse dust foreground can bias estimates of CMB polarization if the chosen dust model cannot adequately represent the true emission spectrum. For the commonly used one- or two-component modified blackbody models, the bias can exceed levels $\Delta r \sim {\rm few} \times 10^{-3}$, large compared to the sensitivities anticipated for the next generation of CMB instrumentation \citep{ kogut/etal:2011, % matsumura/etal:2014, % delabrouille:2015, % abazajian/etal:2015}. % Foreground bias at this level would represent a statistically significant error in derived cosmological parameters and (depending on the cosmology) could induce a false detection of the inflationary signal. Several techniques may be employed to reduce this bias. One method is to restrict the observing bands to frequencies 30 GHz $< \nu <$ 250 GHz where the CMB is brightest compared to the combined synchrotron and dust foregrounds, thereby reducing the precision to which the foreground emission must be modeled. Parametric fitting over a restricted frequency range also reduces the effect of un-modeled spectral curvature present when the fitted spectral model differs from the true foreground spectra. However, the smaller foreground amplitudes within the restricted frequency range require correspondingly better sensitivity in foreground-dominated channels to identify and remove the foreground emission. \begin{figure}[t] \centerline{ \includegraphics[width=3.2in]{show_dust_resid_vs_freq_apj.eps} } \caption{Effect of restricted observing frequencies for parametric dust models. (Top) Bias $\Delta r$ caused by fitting a single-temperature modified blackbody dust model to a sky containing dust emission from the two-temperature, two-Gaussian, or transient heating toy models (see text). The best-fit value for $r$ changes sign when fitting the transient heating input. (Bottom) Residual between the input sky and best-fit model, evaluated at the highest observed frequency. } \label{bias_vs_freq} \end{figure} We quantify this trade by fitting simulated skies while varying the frequency range over which the fit is performed. The input sky consists of power-law synchrotron with spectral index $-1.05$ plus dust specified by either the two-temperature, two-Gaussian, or transient toy models from $\S$4.4. The fitted model includes CMB, synchrotron, and single-temperature modified blackbody dust for a total of 9 free parameters. We fit the simulated sky using 10 observing channels spaced logarithmically in frequency starting at 30 GHz, and vary the highest observed frequency from 150 GHz to 950 GHz to show the effect of extending observations into the brighter dust foreground at higher frequencies. Figure \ref{bias_vs_freq} shows the results. For the two-temperature and two-Gaussian input models, the simulations show the expected pattern with the bias $\Delta r$ monotonically decreasing as the frequency range is restricted. Fitting a single modified blackbody to a toy model with a high-temperature tail shows a different pattern. When the fit is restricted to frequencies below 400 GHz, the single-component model under-estimates the dust emission and over-estimates the CMB. When the frequency range is extended to the higher frequencies where the warmer components dominate, the single-component model over-estimates the dust emission and under-estimates the CMB. The residual CMB term thus changes sign, creating a null for the special case where the frequency coverage extends just high enough that the dust residuals cancel. Note that, for the transient heating toy model, restricting the frequency coverage below 400 GHz actually {\it increases} the bias in $r$. A second consideration for restricted frequency coverage is sensitivity. As the frequency coverage is restricted, the parametric models better reconstruct the superposed emission from the CMB and combined foregrounds. This is not the same as accurately reconstructing the CMB signal itself. For the noiseless simulations in this paper, we define a goodness-of-fit statistic $\Gamma$ using the rms difference between the intensities of the input sky and the best-fit model, summed over observing channels (Eq. \ref{resid_eq}). This fractional residual varies from $\Gamma = 10^{-5}$ for fits restricted to frequencies $\nu < 200$ GHz to $\Gamma = 10^{-3}$ for fits extending to $\nu < 950$ GHz and is nearly identical for all three toy models considered. The toy models evaluated above yield non-trivial bias in $r$ despite fitting the combined sky emission to sub-percent precision. Since the functional form of far-IR dust emission is not known {\it a priori}, distinguishing between biased and unbiased fits requires channel sensitivities comparable to the residuals derived from plausible models of dust emission. The design of CMB missions typically includes one or more ``guard channels'' at higher frequencies where dust emission dominates. The bottom panel of Figure \ref{bias_vs_freq} shows the guard channel sensitivity required to distinguish these toy models, \begin{equation} \delta S = \Gamma ~I_{\rm input} \vert_{\nu = \nu_{\rm max}} , \label{sensitivity_def} \end{equation} evaluated as a function of the highest observing frequency. Since the highest observing frequency may extend to frequencies beyond the Wien cutoff in the CMB intensity, we present the sensitivity $\delta S$ in units of Rayleigh-Jeans temperature, \begin{equation} \delta T_{\rm RJ} = \delta S ~\frac{\lambda^2}{2k} \label{rj_def} \end{equation} where $k$ is Boltzmann's constant and $\lambda$ is the observing wavelength. Figure \ref{bias_vs_freq} illustrates the challenges for observations with restricted frequency coverage. A nine-parameter fit limited to observations at frequencies $\nu < 250$ GHz most accessible to ground-based observations yields biases $\Delta r~>~10^{-3}$ despite fitting the combined sky emission to precision $\Gamma \sim 10^{-5}$. A statistical detection of this bias would require sub-nK channel sensitivities. At least two solutions are possible. Fits with additional free parameters can better model the (unknown) dust spectrum, reducing the bias in the fitted CMB component. A 13-parameter model reduces the bias substantially ($\S$4.4) but would require at least 14 observing channels. Fitting multiple temperature components within a restricted frequency range may additionally lead to parameter degeneracies. An alternative is to include observations at higher frequencies. Extending the observations to THz frequencies where the dust emission peaks provides discrimination among various dust models. For example, extending observations to 3 THz increases the residual between the input sky and best-fit model to 400 Jy sr$^{-1}$ for the TLS model and over 2000 Jy sr$^{-1}$ for the transient heating model. Differences at this level are readily detectable, providing statistical evidence for a poorly fitted model. Observations at frequencies at or above the peak of the dust spectrum can additionally provide information on the actual dust temperature distribution, further discriminating between candidate dust models. Emission from dust components at successively higher temperatures will peak at successively higher frequencies (the well-known Wien displacement law); however, the superposed emission from multiple temperature components will typically show a single broad peak. One indicator of the temperature information within the superposed spectrum may be shown by Gram-Schmidt orthogonalization. We begin with a Planck spectrum at some minimum temperature $T_{\rm min}$, and construct orthogonal linear combinations from Planck spectra at successively higher temperatures to some maximum $T_{\rm max}$. Figure \ref{orthogonal_spectra} shows the resulting orthogonal components for temperatures 8~K $< ~T ~ <$ 24 K, evaluated at frequencies $\nu < 5$ THz. There is relatively little spectral content at frequencies below 800 GHz, but well-separated peaks in the component spectra at higher frequencies. Another way to evaluate the ability of THz observations to constrain the dust temperature distribution is to note that emission on the Rayleigh-Jeans side of the emission peak shows flatter spectral dependence compared to the exponential Wien cutoff at higher frequencies. Dust at temperatures $T_d > 8$ K will peak at frequencies above 800 GHz, so it is not surprising to find most of the spectral content at frequencies above 800 GHz. \begin{figure}[t] \centerline{ \includegraphics[width=2.7in,angle=90]{show_ortho_dust_spectra_apj.eps} } \caption{ Dust spectral components corresponding to different physical temperatures, derived from Gram-Schmidt orthogonalization of the individual Planck spectra at each temperature. The minimal spectral content at frequencies below 800 GHz makes it difficult for observations at low frequencies to distinguish multiple temperature components: most of the spectral content lies at frequencies above 1 THz. } \label{orthogonal_spectra} \end{figure}

1607.02150

1607

1607.00013_arXiv.txt

We present an overview and the first data release of ZFIRE, a spectroscopic redshift survey of star-forming galaxies that utilizes the MOSFIRE instrument on Keck-I to study galaxy properties in rich environments at $1.5<z<2.5$. ZFIRE measures accurate spectroscopic redshifts and basic galaxy properties derived from multiple emission lines. The galaxies are selected from a stellar mass limited sample based on deep near infrared imaging ($\mathrm{K_{AB}<25}$) and precise photometric redshifts from the ZFOURGE and UKIDSS surveys as well as grism redshifts from 3DHST. Between 2013 and 2015 ZFIRE has observed the COSMOS and UDS legacy fields over 13 nights and has obtained 211 galaxy redshifts over $1.57<z<2.66$ from a combination of nebular emission lines (such as \Halpha, \NII, \Hbeta, \OII, \OIII, \SII) observed at 1--2\micron. Based on our medium-band near infra-red photometry, we are able to spectrophotometrically flux calibrate our spectra to \around10\% accuracy. ZFIRE reaches $5\sigma$ emission line flux limits of \around$\mathrm{3\times10^{-18}~erg/s/cm^2}$ with a resolving power of $R=3500$ and reaches masses down to \around10$^{9}$\msol. We confirm that the primary input survey, ZFOURGE, has produced photometric redshifts for star-forming galaxies (including highly attenuated ones) accurate to $\Delta z/(1+z\mathrm{_{spec})}=0.015$ with $0.7\%$ outliers. We measure a slight redshift bias of $<0.001$, and we note that the redshift bias tends to be larger at higher masses. We also examine the role of redshift on the derivation of rest-frame colours and stellar population parameters from SED fitting techniques. The ZFIRE survey extends spectroscopically confirmed $z\sim 2$ samples across a richer range of environments, here we make available the first public release of the data for use by the community.\footnote{\url{http://zfire.swinburne.edu.au}}

The rapid development of very deep multi-wavelength imaging surveys from the ground and space in the past decade has greatly enhanced our understanding of important questions in galaxy evolution particularly through the provision of `photometric redshift' estimates (and hence the evolutionary sequencing of galaxies) from multi-band spectral energy distribution (SED) fitting \citep{Whitaker2011,McCracken2012,Skelton2014}. Studies using data from these surveys have led to a more detailed understanding of topics such as the evolution of the galaxy mass function \cite[eg.,][]{Marchesini2010,Muzzin2013,Tomczak2014,Grazian2015}, stellar population properties \cite[eg.,][]{Maseda2014,Spitler2014,Pacifici2015}, evolution of galaxy morphology \cite[eg.,][]{Huertas-Company2015,Papovich2015}, and the growth of the large-scale structure in the universe \citep{Adelberger2005,Wake2011}. \subsection{Advances with Deep Near-IR Imaging Surveys} Near-infrared data is vital for this endeavour, both for photometric redshift estimation \citep{Dahlen2013,Rafelski2015} and provision of stellar mass estimates \citep{Brinchmann2000,Muzzin2009}. Stellar mass is especially useful for tracking galaxy evolution as it increases monotonically with time, but data at near-infrared wavelengths are needed to estimate it accurately at high-redshift \citep[][Straatman et al. in press]{Whitaker2011}. New surveys have been made possible by the recent development of relatively wide-field sensitive near infrared (NIR) imagers in 4-8m telescopes such as FourStar \citep{Persson2013} , HAWK-I \citep{Pirard2004}, NEWFIRM \citep{NEWFIRM} and VIRCAM \citep{Dalton2006}. Surveys such as ZFOURGE (Straatman et al., in press), the NEWFIRM medium-band Survey (NMBS) \citep{Whitaker2011}, and ULTRAVISTA \citep{McCracken2012} have obtained deep imaging over relatively large sky areas (up to 1.5 deg$^2$). The introduction of near-infrared medium-band filters ($\Delta\lambda\sim 1000$\AA) has resulted in photometric redshifts with accuracies of \around2\% \citep{Whitaker2011} and enabled galaxy properties to be accurately derived by SED fitting techniques such as EAZY \citep{Brammer2008} and FAST \citep{Kriek2009}. These photometric redshift surveys have greatly enhanced our understanding of the universe at $z\sim2$, which is a critical epoch in the evolution of the universe. At this redshift, the universe was only 3 billion years old and was at the peak of cosmic star formation rate activity \citep{Hopkins2006,Lee2015}. We see the presence of massive, often dusty, star-forming galaxies \citep{Spitler2014,Reddy2015} which were undergoing rapid evolution and the development of a significant population of massive, quiescent galaxies \citep{vanDokkum2008,Damjanov2009}. Galaxy clusters have also now been identified at $z\sim2$, and results indicate that this may be the epoch when environment starts to influence galaxy evolution \citep{Gobat2011,Spitler2012,Yuan2014,Casey2015}. \subsection{Need for Spectroscopy} Even though immense progress on understanding galaxy evolution has been made possible by deep imaging surveys, the spectroscopy of galaxies remains critically important. Spectroscopy provides the basic, precision redshift information that can be used to investigate the accuracy of photometric redshifts derived via SED fitting techniques. The galaxy properties derived via photometry have a strong dependence on the redshifts, and quantifying any systematic biases will help constrain the derived galaxy properties and understand associated errors. Spectral emission and absorption lines also provide a wealth of information on physical processes and kinematics within galaxies \citep{Shapley2009}. Spectroscopy also provides accurate environmental information (for example, the velocity dispersions of proto-clusters; e.g. \cite{Yuan2014}) beyond the resolution of photometric redshifts. Rest-frame ultraviolet (UV) spectroscopy of galaxies provides information on the properties of massive stars in galaxies and the composition and kinematics of the galaxies' interstellar medium \citep[ISM;][]{Dessauges2010,Quider2010}. Rest-frame optical absorption lines are vital to determine the older stellar population properties of the galaxies \citep[eg.,][]{vandeSande2011,Belli2014}. Rest-frame optical emission lines provide information on the state of the ionized gas in galaxies, its density, ionization degree, and metallicity \citep{Pettini2004,Steidel2014,Kacprzak2015,Kewley2016,Shimakawa2015}. \subsection{Spectroscopy of $z\lesssim1$ Galaxies} Large-scale spectroscopy is now routine at the low redshift universe. Surveys such as the Sloan Digital Sky Survey \citep[][]{York2000}, the 2-Degree Field Galaxy Redshift Survey \citep[][]{Colless2001}, and the Galaxy and Mass Assembly Survey \citep[][]{Driver2009} extensively explored the $z\la 0.2$ universe ($10^5$--$10^6$ galaxies). At $z\sim 1$ the DEEP2 Galaxy Redshift Survey \citep{Newman2013}, the VIMOS VLT Deep Survey \citep{LeFevre2005}, the VIMOS Public Extragalactic Survey \citep{Garilli2014}, and zCOSMOS \citep{Lilly2007} have produced large spectroscopic samples ($10^4$--$10^5$ galaxies). The large number of galaxies sampled in various environmental and physical conditions by these surveys has placed strong constraints on galaxy models at $z<1$ while revealing rare phases and mechanisms of galaxy evolution \cite[e.g.,][]{Cooper2007,Coil2008,Cheung2012,Newman2013}. \subsection{Spectroscopy of $z\sim2$ Galaxies} At a $z\gtrsim1.5$ rest-frame optical features are redshifted to the NIR regime and therefore accessing these diagnostics becomes more challenging. Historically, the spectroscopy of galaxies in these redshifts focussed on the follow up of Lyman break galaxies, which are rest-frame UV selected using the distribution of the objects in $\cal{U}$, $\cal{G}$, and $\cal{R}$ colour space \citep{Steidel1992}. This technique takes advantage of the discontinuity of the SEDs near the Lyman limit. \citet{Steidel2003} used this technique to target these candidates with multi-object optical spectrographs to obtain rest frame UV spectra for \around1000 galaxies at $z\sim3$. Furthermore, $\cal{U}$, $\cal{G}$, and $\cal{R}$ selections can be modified to select similar star-forming galaxies between $1.5<z<2.5$ via their U-band excess flux \citep{Steidel2004}. Such sample selections are biased toward UV bright sources and do not yield homogeneous mass complete samples. Surveys such as the Gemini Deep Deep Survey \citep[][]{Abraham2004} and the Galaxy Mass Assembly ultra-deep Spectroscopic Survey \citep[][]{Kurk2013} have attempted to address this by using the IR selection of galaxies (hence much closer to mass-complete samples) before obtaining optical spectroscopy. The K20 survey \citep{Cimatti2002} used a selection based on Ks magnitude (Ks$<20$) to obtain optical spectroscopy of extremely dusty galaxies at $z\sim1$. These surveys have provided redshift information, but only rest-frame UV spectral diagnostics, and many red galaxies are extremely faint in the rest-UV requiring very long exposure times. The development of near-IR spectrographs has given us access to rest-frame optical spectroscopy of galaxies at $z\gtrsim1.5$, but the ability to perform spectroscopy of a large number of galaxies has been hindered due to low sensitivity and/or unavailability of multiplexed capabilities. For example the MOIRCS Deep Survey \citep{Kajisawa2006} had to compromise between area, sensitivity, number of targets, and resolution due to instrumental limits with MOIRCS in Subaru \citep{Ichikawa2006}. The Subaru FMOS galaxy redshift survey \cite{Tonegawa2015}, yielded mostly bright line emitters due to limitations in sensitivity of FMOS \citep{Kimura2010}. Furthermore, FMOS does not cover the longer K-band regime which places an upper limit for \Halpha\ detections at $z\sim1.7$. Sensitive long slit spectrographs such as GNIRS \citep{Elias2006} and XShooter \citep{Vernet2011} have been utilised to observe limited samples of massive galaxies at $z\sim2$. NIR-grism surveys from the \emph{Hubble Space Telescope (HST)} have yielded large samples such as in the 3DHST survey \citep{Momcheva2015,Treu2015} but have low spectral resolution ($R\sim70-300$) and do not probe wavelengths $>$ 2\micron. With the introduction of the Multi-object Spectrometer for infrared Exploration (MOSFIRE), a cryogenic configurable multislit system on the 10m Keck telescope \citep{McLean2012}, we are now able to obtain high-quality near-infrared spectra of galaxies in large quantities \citep{Kulas2013,Steidel2014,Kriek2015,Wirth2015}. The Team Keck Redshift Survey 2 observed a sample of 97 galaxies at $z\sim2$ to test the performance of the new instrument \citep{Wirth2015} and investigate the ionization parameters of galaxies at $z\sim2$. The Keck Baryonic Structure Survey is an ongoing survey of galaxies currently with 179 galaxy spectra, which is primarily aimed to investigate the physical processes between baryons in the galaxies and the intergalactic medium \citep{Steidel2014}. The MOSFIRE Deep Evolution Field (MOSDEF) survey is near-infrared selected and aims to observe \around1500 galaxies $1.5<z<3.5$ to study stellar populations, Active Galactic Nuclei, dust, metallicity, and gas physics using nebular emission lines and stellar absorption lines \citep{Kriek2015}. \subsection{The ZFIRE Survey} In this paper, we present the ZFIRE survey, which utilizes MOSFIRE to observe galaxies in rich environments at $z>1.5$ with a complementary sample of field galaxies. A mass/magnitude complete study of rich galaxy environments is essential to overcome selection-bias. Galaxy clusters are the densest galaxy environments in the universe and are formed via various physical processes \citep{Kravtsov2012}. They are a proxy for the original matter density fields of the universe and can be used to constrain fundamental cosmological parameters. Focusing on these rich environments at high-redshift provides access to numerous galaxies with various physical conditions that are rapidly evolving and interacting with their environments. These galaxies can be used to study the formation mechanisms of local galaxy clusters in a period where they are undergoing extreme evolutionary processes. Such environments are rare at $z\sim 2$ \citep{Gobat2011,Newman2014,Yuan2014}: for example, we target the \cite{Spitler2012} cluster at $z=2.1$, which was the only such massive structure found in the 0.1 deg$^2$ ZFOURGE survey (and that at only 4\% chance, \citep{Yuan2014}). Hence, a pointed survey on such clusters and their environs is highly complementary to other field surveys being performed with MOSFIRE. Here we present the ZFIRE\ survey overview and first data release. We release data for two cluster fields: one at $z=2.095$ \citep{Spitler2012,Yuan2014} and the other at $z=1.62$ \citep{Papovich2010,Tanaka2010}. The structure of the paper is as follows: in Section \ref{sec:survey}, we describe the ZFIRE survey design, target selection and data reduction. In Section \ref{sec:results}, we present our data and calculate the completeness and detection limits of the survey. We investigate the accuracy of photometric redshifts of different surveys that cover the ZFIRE\ fields in Section \ref{sec:photometric_redshifts}. In Section \ref{sec:implications}, we study the role of photometric redshift accuracy on galaxy physical parameters derived via common SED fitting techniques and how spectroscopic accuracy affects cluster membership identification. A brief description of the past/present work and the future direction of the survey is presented in Section \ref{sec:summary}. We assume a cosmology with H$_0$= 70 km/s/Mpc, $\Omega_\Lambda$=0.7 and $\Omega_m$= 0.3. Unless explicitly stated we use AB magnitudes throughout the paper. Stellar population model fits assume a \citet{Chabrier2003} initial mass function (IMF), \citet{Calzetti2001} dust law and solar metallicity. We define \zspec\ as the spectroscopic redshift, \zphoto\ as the photometric redshift, and \zgrism\ as the grism redshift from 3DHST \citep{Momcheva2015}. We express stellar mass (M$_*$) in units of solar mass (M$_\odot$). Data analysis was performed using \texttt{iPython} \citep{Perez2007} and \texttt{astropy} \citep{Astropy2013} and \texttt{matplotlib} \citep{Hunter2007} code to reproduce the figures, will be available online\footnote{https://github.com/themiyan/zfire\_survey}.

\label{sec:summary} Here we present the ZFIRE survey of galaxies in rich environments and our first public data release. A detailed description of the data reduction used by ZFIRE is provided. The use of a flux standard star along with photometric data from ZFOURGE and UKIDSS has made it possible to flux calibrate the spectra to $\lesssim10$\% accuracy. The ZFIRE-COSMOS sample spans a wide range in Ks magnitude and stellar mass and secures redshifts for UVJ star-forming galaxies to Ks=24.1 and stellar masses of $\log_{10}($\mass$)>9.3$. We show that selecting using rest-frame UVJ colours is an effective method for identifying \Halpha-emitting galaxies at $z\sim2$ in rich environments. Redshifts have been measured for 232 galaxies of which 87 are identified as members of the rich clusters we have targeted in COSMOS and UDS fields. Photometric redshift probability density functions from EAZY are used to show that the expected \Halpha\ detections are similar to the ZFIRE detection rate in the COSMOS field. In the COSMOS field, the ZFIRE survey has detected \around80\% of the targeted star-forming galaxies. We also show that the density structure discovered by \citet{Spitler2012} has been thoroughly sampled by ZFIRE. Using spectroscopic redshifts from ZFIRE with ZFOURGE and other public photometric survey data, we investigated the accuracies of photometric redshifts. The use of medium-band imaging in SED fitting techniques can result in photometric redshift accuracies of $\sim1.5\%$. ZFIRE calculations of photometric redshift accuracies are consistent with the expectations of the ZFOURGE survey (Straatman at al., in press) but are slightly less accurate than the NMBS \citep{Whitaker2011} and 3DHST \citep{Skelton2014} survey results. The higher redshift errors can be attributed to sampling differences, which arises from the deeper NIR medium-band imaging in ZFOURGE compared to the other surveys (i.e. overlapping galaxies tend to be fainter than typical in the respective galaxies in NMBS). If we select a brighter subset of NMBS (Ks $<23$) we find that the redshift accuracy increases by 30\%. Using UKIDSS, \citet{Quadri2012} shows that the photometric redshift accuracy is dependent on redshift and that at higher redshifts the photometric redshift error is higher. Between UKIDSS at $z\sim1.6$ and ZFOURGE at $z\sim2$ the photometric redshift accuracies are similar. Therefore, the use of medium-band imaging in ZFOURGE has resulted in more accurate redshifts at $z\sim2$, due to finer sampling of the D4000 spectral feature by the J1, J2, and J3 NIR medium-band filters. The introduction of medium-bands in the K band in future surveys may allow photometric redshifts to be determined to higher accuracies at $z\gtrsim4$. The importance of spectroscopic surveys to probe the large-scale structure of the universe is very clear. For the COSMOS \citet{Yuan2014} cluster, we compute a 38\% success rate (i.e., 38\% of galaxies in $3\sigma$ overdensity regions are identified spectroscopically as cluster galaxies) and a 56\% incompleteness (56\% of spectroscopic cluster galaxies are not identified from data based on purely photometry) using the best photometric redshifts (with seventh nearest neighbour algorithms) to identify clustered galaxies. We find a systematic trend in photometric redshift accuracy, where massive galaxies give higher positive offsets up to $\sim$0.05 for $\Delta z/(1+z_\mathrm{spec}$) values as a function of galaxy stellar mass. However, it is not evident that there is any statistically significant trend for a similar relationship with galaxy luminosity. Results also suggest that the stellar mass and SFR correlates with redshift error. This is driven by the change in the calculated galaxy luminosity as a function of the assigned redshift and we show that the values correlate approximately with the theoretical expectation. SFR shows larger scatter compared to stellar mass in this parameter space, which can be attributed to the stronger weight given to UV flux, which is very sensitive to the underlying model, in the derivation of the SFR. This stronger correlation of the UV flux with redshift error is further evident when comparing the change in (U$-$V) and (V$-$J) colour with change in redshift. When rest-frame U,V, and J colours are re-derived using spectroscopic redshifts, our results show a stronger change in (U$-$V) colour compared to the (V$-$J) colour. Therefore, a redshift error may introduce an extra selection bias on rest-frame UVJ selected galaxies. Further studies using larger samples of quiescent and dusty star-forming galaxies at $z\sim2$ are needed to quantify this bias. Clearly the use of photometric redshifts can lead to biases even when using the same SED template set. However, it is important to acknowledge the underlying uncertainties that lie in deriving galaxy properties even with spectroscopic redshifts. Future work could consider the role of SED templates used in SED fitting techniques. Generally the templates used are empirically derived, which limits the capability to understand the inherent properties of the observed galaxies. With the use of physically motivated models such as MAGPHYS \citep{daCunha2008}, more statistically meaningful relationships between different physical parameters of the observed galaxies could be obtained. Improving such models to include photo-ionization of galaxies the in future will allow us to directly make comparisons of star-forming galaxies at $z\sim2$, which will be vital to study the inherent galaxy properties. Furthermore, the accuracy of underlying assumptions used in SED fitting techniques such as the IMF, dust properties, and star formation histories at $z\sim2$ should be investigated. These assumptions are largely driven by observed relationships at $z\sim0$, and if the galaxies at higher redshifts are proven to be inherently different from the local populations, results obtained via current SED fitting techniques may be inaccurate. Future work should focus on the physical understanding of the galaxy properties at $z\gtrsim2$ with large spectroscopic surveys to better constrain the galaxy evolution models. The recent development of sensitive NIR integral field spectrographs with multiplexed capabilities will undoubtedly continue to add a wealth of more information on this topic over the next few years. The ZFIRE survey will continue focusing on exploring the large spectroscopic sample of galaxies in rich environments at $1<z<3$ to investigate galaxy properties in rich environments. Upcoming papers include analyses of the IMF (T. Nanayakkara et al. 2016, in preparation), kinematic scaling relations (\citet{Alcorn2016}; C. Straatman et al. 2016, in preparation), the mass--metallicity fundamental plane \citep{Kacprzak2016}, and galaxy growth in cluster and field samples (K. Tran et al., in preparation).

1607.00013

1607

1607.03427_arXiv.txt

We report results of the monitoring campaign of the transient X-ray pulsar \smc\ performed with the {\it Swift}/XRT telescope in the period of Sept 2015 -- Jan 2016 during the Type II outburst. During this event bolometric luminosity of the source ranged from $\simeq10^{39}$ down to $several \times 10^{34}$~\lum. Moreover, we discovered its dramatic drop by a factor of more than 100 below the limiting value of $L_{\rm lim}\simeq4\times10^{36}$~\lum, that can be interpreted as a transition to the propeller regime. These measurements make \smc\ the sixth pulsating X-ray source where such a transition is observed and allow us to estimate the magnetic field of the neutron star in the system $B\simeq3\times10^{12}$~G, that is in agreement with independent results of the spectral analysis.

\label{sec:intro} The Small Magellanic Cloud (SMC) is the satellite of the Milky Way situated at the distance of $d\simeq62$~kpc \citep{2012AJ....144..107H}. This galaxy is extremely rich in Be X-ray binary systems, harboring a neutron star orbiting around an OBe companion \citep[see recent review of ][and references therein]{2015MNRAS.452..969C}. \smc\ was discovered at an early stage of the study of SMC \citep{1978ApJ...221L..37C} during the Type II outburst with an X-ray luminosity in the 2--11 keV energy band of about $10^{38}$\,\lum. A pulsating nature of the source was established during the second registered outburst, when the pulsations with the period of $P_{\rm spin}\simeq2.37$ s were detected with the {\it RXTE} observatory from the sky region around \smc\ \citep{2001ApJ...548L..41C} and confirmed later using the {\it ASCA} data \citep{2001PASJ...53..227Y}. The pulsar magnetic field was not known until present time, but some hint at the cyclotron line detection in the \smc\ spectrum near $\sim27$ keV was reported recently by \citet{cycline}. The optical counterpart of \smc\ was not unambiguously identified during a quite long time as two different stars of early spectral type are located near the X-ray position of the source with the angular separation only 2.5\arcsec. Only recent monitoring observations with the OGLE experiment allowed to reveal a variability of one of these stars with the period of $P_{\rm orb}=18.62\pm0.02$ days \citep{2011MNRAS.412..391S}, that is in agreement with periodical variations of the pulse period detected by the {\it RXTE} and {\it Swift} observatories at $P_{\rm orb}\simeq18.4$ days \citep{2011MNRAS.416.1556T, 2016MNRAS.458L..74L}. This periodicity was interpreted as an orbital period in the system that, in a combination with the pulse period of $P_{\rm spin}\simeq2.37$~s, places \smc\ in the Be-systems region in the Corbet diagram \citep{1986MNRAS.220.1047C}. The transient nature of a majority of Be systems is ideally suited to study the magnetic field, to investigate luminosity-dependent accretion processes and the geometry of the system, as well as to learn about the interaction of the accreted matter with the neutron star magnetosphere \citep[see, e.g., ][for reviews and current physical models]{1998A&A...338..505N, 2011Ap&SS.332....1R, 2013ApJ...777..115P,walter15}. One of the most interesting and straightforward manifestation of such an interaction is a transition of the accreting neutron star to the so-called propeller regime. The physical aspects of this regime was considered by \citet{1975A&A....39..185I}, who showed that under some conditions the accreted matter can be stopped by the centrifugal barrier set up by the rapidly rotating magnetosphere of the strongly magnetized neutron star. It should lead to the dramatic drop of the X-ray intensity of the source. The moment of the transition from the normal accretion regime to the propeller one depends on a combination of three physical parameters of the system -- the pulse period, magnetic moment (or magnetic field strength) of the neutron star and the accretion rate. Because the pulse period and the accretion rate can be derived from observations, the detection of the propeller effect provides us with an independent estimation of the neutron star magnetic field. This knowledge is very important as the magnetic field is one of the fundamental parameters governing observed properties of neutron stars. Until recently only a few cases of possible transitions into the propeller regime in accreting millisecond and X-ray pulsars were reported in the literature \citep{1986ApJ...308..669S,1997ApJ...482L.163C,2001ApJ...561..924C,2008ApJ...684L..99C}. Recently the propeller effect was also discovered in the first pulsating ultra-luminous X-ray source M82\,X-2 \citep{2016MNRAS.457.1101T}. This discovery initiated a special monitoring program of transient X-ray pulsars with the {\it Swift}/XRT telescope to search for the propeller effect in other sources. First results of this program were published by \cite{2016A&A...593A..16T} for two well known transient X-ray pulsars V\,0332+53 and 4U\,0115+634, where the propeller effect was firmly established. In this paper we report a discovery of the propeller effect and consequent determination of the magnetic field strength in another transient X-ray pulsar \smc.

In this paper we have reported the discovery of the propeller effect in the bright transient X-ray pulsar \smc. The dramatic drop of the source luminosity (by a factor of more than 100) on the time scale of a few days was revealed thanks to the monitoring campaign with the {\it Swift}/XRT telescope, organized during the Type II outburst registered from the source in 2015 September -- 2016 January. The luminosity drop occurred near the luminosity of $L_{\rm lim}\simeq4\times10^{36}$~\lum. Based on this measurement we estimated the magnetic field strength of the neutron star in the \smc\ binary system as $B\simeq3\times10^{12}$~G, that is typical for X-ray pulsars \citep[see, e.g., recent review by][]{walter15} and confirmed independently by results of the spectral analysis. Thus our discovery makes \smc\ the sixth known pulsating X-ray source where the propeller effect was observed.

1607.03427

1607

1607.08818_arXiv.txt

A spectral line image cube generated from 115 minutes of MWA data that covers a field of view of 400 sq. deg. around the Galactic Centre is used to perform the first Search for ExtraTerrestrial Intelligence (SETI) with the Murchison Widefield Array. Our work constitutes the first modern SETI experiment at low radio frequencies, here between 103 and 133 MHz, paving the way for large-scale searches with the MWA and, in the future, the low frequency Square Kilometre Array. Limits of a few hundred mJy/beam for narrow band emission (10 kHz) are derived from our data, across our 400 sq. deg. field of view. Within this field, 45 exoplanets in 38 planetary systems are known. We extract spectra at the locations of these systems from our image cube, to place limits on the presence of narrow line emission from these systems. We then derive minimum isotropic transmitter powers for these exoplanets; a small handful of the closest objects (10s of pc) yield our best limits of order $10^{14}$ W (Equivalent Isotropic Radiated Power: EIRP). These limits lie above the highest power directional transmitters near these frequencies currently operational on Earth. A SETI experiment with the MWA covering the full accessible sky and its full frequency range would require approximately one month of observing time. The MWA frequency range, its Southern Hemisphere location on an extraordinarily radio quiet site, its very large field of view, and its high sensitivity make it a unique facility for SETI.

The first modern Search for ExtraTerrestrial Intelligence (SETI) was undertaken at radio wavelengths in 1960 at Green Bank, West Virginia \citep{dra08}, targeting two stars, Tau Ceti and Epsilon Eridani at frequencies near the 21-cm line of neutral hydrogen. In the decades since, SETI programs have continued to be undertaken at radio wavelengths, on the basis that highly sensitive radio telescopes exist for astronomy, the radio band is the cornerstone of communications technologies on Earth, and it could be reasonably assumed that a similar technology path has been taken by extraterrestrial civilisations. In 1985, the Million-channel ExtraTerrestrial Assay (META) was established at Harvard University \citep{lei97}, also near 1.4 GHz. META was upgraded to the Billion-channel ExtraTerrestrial Assay (BETA) in 1995, in the extended frequency range of 1.4 - 1.7 GHz. The SERENDIP (Search for Extraterrestrial Radio Emissions from Nearby Developed Intelligent Populations) program was established in 1978 and has evolved considerably since then \citep{wer01}. The SERENDIP program also gave rise to projects such as SETI@home \citep{kor09} and Southern SERENDIP \citep{sto00}, a survey conducted using the Parkes radio telescope. A novel targeted search using Very Long Baseline Interferometry observations of the Gliese 581 star system is described by \citet{ram12}. The Allen Telescope Array (ATA) has been used extensively for a range of SETI experiments over the last decade \citep{wel09}. See \citet{gar14} for a recent review of aspects of SETI experiments at radio wavelengths. For general reviews, see \citet{sie15,tar03}. The SETI experiments conducted at radio wavelengths to date have generally focused on the 1.4 - 1.7 GHz range, the so-called ``water hole" between the prominent radio spectral lines due to neutral hydrogen (H) and hydroxyl (OH). However, many other radio frequencies are also viable for SETI experiments. One radio frequency range that has opened up in recent years is in the tens to hundreds of MHz. Powerful multi-purpose, next-generation low frequency radio telescopes such as (LOFAR: \citealt{van14}) and the Murchison Widefield Array (MWA: \citealt{tin13,lon09}) are precursors and pathfinders for the billion dollar Square Kilometre Array (SKA: \citealt{dew09}) over the next decade. A key science program for the SKA is the ``cradle of life" \citep{hoa15}, including comprehensive and ambitious SETI experiments \citep{sie15}. \citet{loe07} discuss the prospects for SETI experiments using facilities such as the MWA, LOFAR, and the low frequency SKA, instruments that are primarily designed to search for the redshifted neutral hydrogen line from the Epoch of Reionisation. \citet{rah15} discusses the potential for utilising microlensing for SETI programs at radio wavelengths, with particular reference to the MWA and SKA. LOFAR, operating in the ranges of 30 - 80 MHz and 120 - 240 MHz, launched a SETI project in 2010\footnote{https://www.astron.nl/about-astron/press-public/news/lofar-opens-low-frequency-universe-and-starts-new-seti-search/lofar-o}, although no results have thus far been published. The MWA operates in the frequency range of 80 - 300 MHz on an exceptionally radio quiet site in Western Australia. The high sensitivity of the MWA, its radio quiet location \citep{2015PASA...32....8O}, its frequency range, its access to the Southern Hemisphere, and its exceptionally large field of view (hundreds of square degrees) make it a unique facility for exploratory SETI experiments. As \citet{gar14} points out, the emergence of new radio telescopes with very large fields of view opens up new areas of parameter space for SETI experiments. For example, within a single $>$500 sq. deg. field of view typical for the MWA, on average tens of stellar systems within 50 lightyears will be accessible \citep{hip97}. Based on recent results showing that planets are the norm rather than the exception, for example on average $1.0\pm0.1$ planets per M dwarf star in our Galaxy \citep{swi13}, one would therefore expect dozens of nearby (within 50 lightyears) planets in a single MWA field of view and far greater numbers of more distant planets. The MWA field of view therefore results in a significant multiplex advantage that can be exploited for SETI experiments. While facilities such as the MWA, LOFAR, and the SKA will open up unprecedented parameter space for new SETI experiments, it is worth noting that the very first consideration of low radio frequency SETI came approximately 60 years before the first modern experiments described in \citet{dra08}, at the dawn of radio communications. In the late 1800s and early 1900s, Guglielmo Marconi\footnote{Reported extensively in the media at the time, for example on page 3 of the New York Tribune, September 2, 1921: http:\/\/chroniclingamerica.loc.gov\/lccn/sn83030214\/1921-09-02\/ed-1\/seq-3\/} and Nicola Tesla \footnote{As described by Corum \& Corum (1996): http:\/\/www.teslasociety.com\/mars.pdf} believed that radio waves could be used to communicate with civilisations on Mars (the widespread belief in the existence of Martian canals persisted at the time) and both claimed to have detected potential signals from that planet. In 1924, an experiment to listen for signals was organised by the US Navy during the opposition of Mars that year, coordinated with a planned cessation of terrestrial radio broadcasts\footnote{http:\/\/www.lettersofnote.com\/2009\/11\/prepare-for-contact.html}. These very first low radio frequency experiments returned null results. In this Letter, we present a first, and opportunistic, SETI pilot experiment with the MWA, in the frequency range 103 - 133 MHz, placing limits on narrow band radio emission toward 38 known planetary systems. The experiment is opportunistic in the sense that the observations were undertaken for a spectral line survey of the Galactic Plane that is ongoing; utility of the data for a SETI experiment was realised post-observation. We use this pilot study to motivate a deeper and significantly larger SETI experiment with the MWA, that could use the full 80 - 300 MHz frequency range and survey the entire southern sky (majority of the Milky Way) visible from Western Australia. The field of SETI has recently received a substantial boost, with the ``Breakthrough Listen" project recently initiated\footnote{http://www.breakthroughinitiatives.org/}. Novel and diverse SETI experiments that sit on the path to utilisation of the SKA for SETI, such as described here, are likely to be useful contributors to such initiatives.

The great majority of exoplanet systems listed in Table 1 are at large distances, yielded from microlensing experiments toward the Galactic Centre, meaning that the observational limits on detectable transmitter power from the MWA are very high, inferred isotropic powers of $10^{17} - 10^{19}$ W. Even if a directional transmitting antenna is assumed, with a gain similar to low frequency over-the-horizon radar transmitters on Earth, the limits on transmitter power are only reduced by factors of order 100. A small handful of exoplanet systems in Table 1 are close enough that the inferred isotropic transmitter powers are of order $10^{13} - 10^{14}$ W. These are still very large in terms of transmitters on Earth. The highest power low frequency transmitters on Earth are the over-the-horizon (OTH) backscatter radars used for military surveillance; these typically operate in the 5 - 30 MHz range and have transmitter powers of order 1 MW. For example, the Jindalee Operation Radar Network (JORN) in Australia has a transmitter power of 560 kW \citep{col00} and similar installations in the US, such as the in the AN/FPS-118 OTH-Backscatter radar\footnote{http://www.globalsecurity.org/wmd/systems/an-fps-118.htm}, have transmitter powers of 1 MW (but can range up to 10 MW). In the latter case of the US system, the Effective Radiated Power is 100 MW, still a factor of $\sim10^{5}$ below the limits for the nearest exoplanet systems in Table 1. Even the addition of the signals from the ensemble of global array of OTH radars fall well below our limits. \cite{loe07} summarise other Earth-based transmitter characteristics relevant to low frequency telescopes. The most powerful transmission ever broadcast deliberately into space was the Arecibo message, directed as a purposeful communication at the globular cluster M13, in a 10 Hz bandwidth at 2380 MHz \citep{ca75}. This transmission had an equivalent isotropic transmission of $20\times10^{9}$ W. Taking the narrow bandwidth into account (and ignoring the large difference in frequency), this transmission once again falls below the limits calculated in Table 1. In the MWA frequency range, \citet{mck13} previously estimated the Equivalent Isotropic Power of FM radio transmissions from the Earth to be 77 MW. This estimate was made by measuring the amount of stray FM radio signal reflected off the Earth's Moon. Again, this isotropic power is well below the limits in Table 1. These projections of Earth-based technologies of course discount the possibility that higher power and/or more highly directive antenna technologies are utilised by advanced extraterrestrial civilisations for communications or remote sensing applications. While the inferred transmitter powers in Table 1 are high compared to the most powerful low frequency transmitters on Earth, this study has nonetheless provided the most comprehensive search for narrow band transmissions from exoplanets in this frequency range. Due to the southern, RFI free location of the MWA, its operational frequency range, and its wide field of view, the MWA provides a unique capability for future SETI projects. This experiment examined one field of view of 400 sq. deg. in a 30.72 MHz frequency band. To perform a SETI experiment to the same depth as achieved here, but over the full MWA frequency range (80 - 300 MHz), and over the full accessible sky from Western Australia, would require approximately one month of observing time. This is an entirely feasible goal for the near future (three times deeper would require of order a year of observing). Moreover, to relieve the restriction of 10 kHz frequency resolution present in the current experiment, it is possible to record voltage data from the MWA and reconstruct coherent beams at far higher spectral resolution to target individual exoplanet systems \cite{tre14}. For example, generating 1 Hz channels from coherent beams across the full array would yield a factor of 100 improvement in sensitivity (assuming a 1 Hz transmission bandwidth), compared to the current experiment. Such a mode could be run communally with the large-scale survey described above, for a selected list of target systems. The current experiment and the capabilities of the MWA provide a clear path to the far greater capabilities of the low frequency component of the Square Kilometre Array, which will be built at the same location as the MWA and have a spectral sensitivity some tens of times greater than the MWA. The radio quiet nature of the MWA/SKA site in enabling SETI experiments at low frequencies (especially through the FM band), as demonstrated here, bodes well for SETI experiments with the SKA.

1607.08818

1607

1607.06458_arXiv.txt

We use the HII galaxies \lsigG\ relation and the resulting Hubble expansion cosmological probe of a sample of just 25 high-$z$ (up to $z \sim 2.33$) \hii\ galaxies, in a joint likelihood analysis with other well tested cosmological probes (CMB, BAOs) in an attempt to constrain the dark energy equation of state (EoS). The constraints, although still weak, are in excellent agreement with those of a similar joint analysis using the well established SNIa Hubble expansion probe. Interestingly, even with the current small number of available high redshift HII galaxies, the HII/BAO/CMB joint analysis gives a 13\% improvement of the quintessence dark energy cosmological constraints compared to the BAO/CMB joint analysis. We have further performed extensive Monte Carlo simulations, with a realistic redshift sampling, to explore the extent to which the use of the \lsigG\ relation, observed in \hii\ galaxies, can constrain effectively the parameter space of the dark energy EoS. The simulations predict substantial improvement in the constraints when increasing the sample of high-$z$ \hii\ galaxies to 500, a goal that can be achieved in reasonable observing times with existing large telescopes and state-of-the-art instrumentation.

The observational evidence for an accelerated cosmic expansion was first given by Type Ia Supernovae (SNIa) \citep{Riess1998, Perlmutter1999}. Since then, measurements of the cosmic microwave background (CMB) anisotropies \citep[e.g.][]{Jaffe2001, Pryke2002, Spergel2007, Planck2015} and of Baryon Acoustic Oscillations (BAOs) \citep[e.g.][]{Eisenstein2005}, in combination with independent Hubble parameter measurements \citep[e.g.][]{Freedman2012}, have provided ample evidence of the presence of a dark energy (DE) component in the Universe. To the present day, the main geometrical tracer of the cosmic acceleration has been SNIa at redshifts $z \lesssim 1.5$ \citep[e.g.][]{Suzuki2012, Betoule2014}. It is of great importance to use alternative geometrical probes at higher redshifts in order to verify the SNIa results and to obtain more stringent constrains in the cosmological parameters solution space \citep{Plionis2011}, with the final aim of discriminating among the various theoretical alternatives that attempt to explain the accelerated expansion of the Universe \citep[cf.][]{Suyu2012}. The \lsig\ relation between the velocity dispersion ($\sigma$) and Balmer-line luminosity ($L[\mathrm{H}x]$, usually H$\beta$) of \hii\ galaxies has already proven its potential as a cosmological tracer \citep[e.g.][and references therein]{Melnick2000, Siegel2005, Plionis2011, Chavez2012, Chavez2014, Terlevich2015}. It has been shown that the \lsigb\ relation can be used in the local Universe to constrain the value of $H_0$ \citep{Chavez2012}. At high-$z$ it can set constraints on the parameters of the DE Equation of State (EoS) \citep{Terlevich2015}. \hii\ Galaxies are a promising tracer for the parameters of the DE EoS precisely because they can be observed, using the current available infrared instrumentation, up to $z \sim 3.5$ \citep[cf.][]{Terlevich2015}. Even when their scatter on the Hubble diagram is about a factor of two larger than in the case of SNIa, this disadvantage is compensated by the fact that \hii\ galaxies are observed to much larger redshifts than SNIa where the degeneracies for different DE models are substantially reduced \citep[cf.][]{Plionis2011}. In addition, because the \lsigb\ relation systematic uncertainty sources\citep{Chavez2012, Chavez2014} are not the same as those of SNIa, \hii\ galaxies constitute an important complement to SNIa in the local Universe, contributing to a better understanding of the systematic errors of both empirical methods. In this paper we perform an HII/BAO/CMB joint likelihood analysis and compare the resulting cosmological constraints with those of a BAO/CMB and a SNIa/BAO/CMB joint likelihood analysis (for the latter we use the \emph{Union 2.1} SNIa compilation \citep{Suzuki2012}). Furthermore, we present extensive Monte-Carlo simulations, tailored to the specific uncertainties of the \hii\ galaxies \lsigb\ relation and currently available instrumentation, to demonstrate its potential possibilities as a cosmological tracer to $z\lesssim 3.5$, to probe a region where the Hubble function is very sensitive to the variations of cosmological parameters \citep{Melnick2000, Plionis2011}. The paper is organised as follows: in section 2 we succinctly describe the data used and associated systematic uncertainties; cosmological constraints that can be obtained from the data are explored in section 3; in section 4 we discuss the Monte-Carlo simulations, in section 5 we discuss the planned data acquisition in order to obtain better constraints on the cosmological parameters. Finally in section 6 we present our conclusions.

\begin{itemize} \item The FoM of the QDE EoS constraints, provided by the joint HII/BAO/CMB analysis, was found to be larger by 13 percent than those provided by the BAO/CMB joint analysis, even with the very small sample of only 25 high-z HII galaxies. \item Both the QDE and CPL EoS constraints of the HII/BAO/CMB and of the SNIa/BAO/CMB joint analyses are in excellent consistency with each other, although (as expected) the SNIa probe still provides a significantly larger FoM. \end{itemize} We have also performed Monte-Carlo simulations tailored to the specific uncertainties of the \lsigb\ relation and to the technical instrumental requirements of KMOS/VLT (and instruments like it). They address the important question of what is the expected increase of the FoM as a function of the number of high-z \hii\ galaxies in the redshift windows accessible. Our previous simulations \citep[cf.][]{Plionis2011} did not take into account the specific error budget of our \lsigb\ relation, or the characteristics of the instruments available and of the accessible redshifts. We would like to add that cosmological analyses, like the one presented in this work, demands a thorough understanding of the interplay between observational random and systematic errors and biases, for which mock catalogues are an essential tool.

1607.06458

1607

1607.06172_arXiv.txt

The Primordial Inflation Polarization ExploreR (PIPER) is a balloon-borne telescope designed to measure the polarization of the Cosmic Microwave Background on large angular scales. PIPER will map 85\% of the sky at 200, 270, 350, and\,600 GHz over a series of 8 conventional balloon flights from the northern and southern hemispheres. The first science flight will use two $32\times40$ arrays of backshort-under-grid transition edge sensors, multiplexed in the time domain, and maintained at 100\,mK by a Continuous Adiabatic Demagnetization Refrigerator. Front-end cryogenic Variable-delay Polarization Modulators provide systematic control by rotating linear to circular polarization at 3 Hz. Twin telescopes allow PIPER to measure Stokes $I$, $Q$, $U$, and $V$ simultaneously. The telescope is maintained at 1.5\,K in an LHe bucket dewar. Cold optics and the lack of a warm window permit sensitivity at the sky-background limit. The ultimate science target is a limit on the tensor-to-scalar ratio of $r\sim0.007$, from the reionization bump to $l\sim300$. PIPER's first flight will be from the Northern hemisphere, and overlap with the CLASS survey at lower frequencies. We describe the current status of the PIPER instrument.

\label{sec:intro} \begin{figure} [ht] \begin{center} \includegraphics[height=6cm]{Figure1.png} \end{center} \caption[power spectrum] { \label{fig:piperplaincl} Left: The polarization of the CMB can be decomposed into E- and B-modes. Right: B-mode power spectrum assuming a tensor-to-scalar ratio $r=0.02$ from inflationary gravitational waves (black) and gravitational lensing (blue). The quantity being measured, the B-mode power, is plotted against angular scale (large multipole $l$ = small scales). PIPER is the only experiment with the sensitivity to measure both the B-mode signal and the dust foreground at large angular scales, where the inflationary signal is not contaminated by the signal from lensing.} \end{figure} Cosmological inflation is the current leading explanation for the observed characteristics of the universe. This theory postulates a period of exponential expansion just after the Big Bang (at $t\approx10^{-35}$ seconds), when the universe expanded by at least 60 e-folds. While the existence of an inflationary epoch can account for observations, no direct evidence yet exists that inflation occurred. The simplest models of inflation predict that the rapid expansion would have produced tensor perturbations in the space-time metric of the universe, i.e. gravitational waves. The current best method for detecting these gravitational waves is to look for the imprint they would have left in the Cosmic Microwave Background (CMB). This imprint comes in the form of a specific pattern, known as B-modes, in the polarization of the CMB. Figure \ref{fig:piperplaincl} shows the polarization signature that CMB experiments are seeking. The observed pattern of linear polarization can be decomposed into curl-free and curl components, known as E-modes and B-modes, respectively, named in analogy to electric and magnetic fields. While density fluctuations can only produce E-modes, gravitational waves can produce both E-modes and B-modes. Hence a detection of B-modes is considered the ``smoking gun" for inflation. The level of the B-mode signal, parameterized by the ratio, $r$, of tensor-to-scalar perturbations, is directly related to the energy scale of inflation. The simplest models of inflation predict $r \sim 0.01$, associating it with the physics of Grand Unified Theories at $10^{16}$\,GeV. A B-mode signal would be a direct probe of the universe at energy scales well beyond the reach of high energy particle colliders. At this level, this signal would be detectable by modern CMB experiments. Discovering the signature of inflation in the polarization of the CMB would advance our understanding of fundamental physics and of the origin of our universe. An inflationary B-mode power spectrum has a characteristic shape, with a ``reionization bump" at large angular scales ($>20^{\circ}$) and a ``recombination peak" at a scale $\sim2^{\circ}$ (Figure \ref{fig:piperplaincl}). The primordial B-mode signal from recombination becomes progressively more contaminated by gravitational lensing of E-modes into B-modes on angular scales smaller than several degrees. The signal is further complicated by the polarized emission of our own galaxy, in particular from interstellar dust and synchrotron radiation. The inflationary B-mode signal is small, at the level of $10^{-9}$\,K, and therefore detecting it requires a high degree of control over systematics and careful removal of galactic foregrounds. \begin{figure} [ht] \begin{center} \includegraphics[height=8cm]{Figure2.pdf} \end{center} \caption[instrument] { \label{fig:instrument} Left: Model of the PIPER payload. Right: PIPER telescope frame. PIPER's twin telescopes are mounted on a fully stainless steel frame in an open-aperture 3500 liter LHe bucket dewar. The telescope is registered to a rigid backbone, and the entire backbone is mounted into the dewar. Only elements of one telescope are labeled; the two telescopes are sagittal mirrors of one another. The first optical element is the VPM, ensuring instrumental polarization cannot be modulated and become systematics. A retractable lid cover protects the detectors on the ground and from the sun. The payload moves in azimuth using a pair of control moment gyros. A star camera provides absolute pointing information.} \end{figure} Joint analysis \cite{Keck2016} of the data from BICEP2/Keck and \textit{Planck} experiments yields an upper limit of $r<0.07$. However this limit relies heavily on modeling of the foreground emission. Two critical features are needed for any experiment aiming to distinguish an inflationary signal from competing foregrounds. One is the ability to probe large angular scales, where the B-mode signal from the reionization bump does not suffer any contamination from lensing. The other is the ability to measure the dominant dust foreground signal with sensitivity better than \textit{Planck}. PIPER is currently the only instrument capable of doing both, by mapping a large fraction of the sky in four bands, at 200, 270, 350 and 600\,GHz, providing the necessary data to measure the large scale B-modes as well as the dust spectrum. By flying each frequency of PIPER in both the Southern and Northern Hemisphere, PIPER will be able to map 85\% of the sky. This will allow the shape of the B-mode power spectrum to be measured at angular scales up to $90^{\circ}$, constraining the characteristic reionization bump. Figure \ref{fig:piperplaincl} shows the sensitivity of PIPER compared to an inflationary model with $r=0.02$. PIPER will also measure the polarized dust foreground to a signal-to-noise of better than 10 even for regions of low dust intensity and even for polarization fractions of 10\%. The PIPER design allows a clean measurement of CMB polarization over the full sky, without the need for complicated scan strategies or boresight rotation.

PIPER will map the polarization of the CMB in four frequency bands (200, 270, 350, and 600\,GHz) over 85\% of the sky over a series of conventional balloon flights from the Northern and Southern Hemisphere. This will allow the shape of the B-mode power spectrum to be measured at angular scales up to $90^{\circ}$, constraining the characteristic reionization bump. After eight flights PIPER will constrain the tensor-to-scalar ratio to $r<0.007$. PIPER will also provide a high signal-to-noise map of the polarized emission from galactic dust, producing a high fidelity template that can be used by CMB experiments to remove the foreground signal from their measurements.

1607.06172

1607

1607.06491_arXiv.txt

We present a newly developed time-dependent three-dimensional multi-zone hadronic blazar emission model. By coupling a Fokker-Planck based lepto-hadronic particle evolution code 3DHad with a polarization-dependent radiation transfer code, 3DPol, we are able to study the time-dependent radiation and polarization signatures of a hadronic blazar model for the first time. Our current code is limited to parameter regimes in which the hadronic $\gamma$-ray output is dominated by proton synchrotron emission, neglecting pion production. Our results demonstrate that the time-dependent flux and polarization signatures are generally dominated by the relation between the synchrotron cooling and the light crossing time scale, which is largely independent of the exact model parameters. We find that unlike the low-energy polarization signatures, which can vary rapidly in time, the high-energy polarization signatures appear stable. As a result, future high-energy polarimeters may be able to distinguish such signatures from the lower and more rapidly variable polarization signatures expected in leptonic models.

Blazars are the most violent class of active galactic nuclei. Their emission is known to be nonthermal-dominated, covering the entire electromagnetic spectrum from radio up to TeV $\gamma$-rays, with strong variability on all time scales \citep[e.g.,][]{Aharonian07}. Blazar spectral energy distributions (SEDs) are characterized by two broad, non-thermal components. The low-energy component, from radio to optical-UV, is generally agreed to be synchrotron radiation of ultrarelativistic electrons. The origin of the high-energy component, from X-rays to $\gamma$-rays, is still under debate. The leptonic model argues that the high-energy component is due to the inverse Compton scattering of either the low-energy synchrotron emission \citep[SSC, e.g.][]{Marscher85,Maraschi92} or external photon fields \citep[EC, e.g.,][]{Dermer92,Sikora94}, while the hadronic model suggests that the high-energy emission is dominated by synchrotron emission of ultrarelativistic protons and the cascading secondary particles resulting from photo-pion and photo-pair production processes \citep[e.g.,][]{Mannheim92,Mucke01}. It is of high importance to many aspects of high energy astrophysics to distinguish these two models, because it will put strong constraints on the blazar jet power, the physics of the central black hole, the origin of ultra-high-energy (UHE) cosmic rays and very-high-energy (VHE, i.e., TeV -- PeV) neutrinos. However, both models are generally able to produce reasonable fits to snap-shot SEDs of blazars \citep[e.g.,][]{Boettcher13}. Thus, additional diagnostics are necessary. An obvious choice would be through the identification of blazars as the sources of VHE neutrinos, which are the ``smoking gun'' of hadronic interactions \citep[e.g.,][]{Halzen97,Kistler14,Diltz15,Petropoulou15}. IceCube has reported detection of astrophysical VHE neutrinos, and there are hints that the origin of these neutrinos could be spatially connected to blazars \citep[e.g.,][]{Aartsen13,Kadler16}. However, in view of the low angular resolution of IceCube, so far the sources of these neutrinos are still unknown. An alternative is the study of light curves. The development of time-dependent leptonic models has been quite fruitful \citep[e.g.,][]{Joshi11,Diltz14,Weidinger15,Asano15}. Although one-zone leptonic models sometimes have difficulty in explaining the frequently seen symmetric light curves, some multi-zone leptonic models that explicitly include the light travel time effects (LTTEs) have successfully resolved that issue \citep[e.g.,][]{Chen14}. On the other hand, due to the more complicated cascading processes, hadronic models are generally stationary and/or single-zone \citep[e.g.,][]{Mastichiadis95,Cerruti15,Yan15}. Another possible discriminant is that leptonic and hadronic models require very distinct magnetic field conditions. Radio to optical polarization measurements have been a standard probe of the jet magnetic field. In particular, recent observations of $\gamma$-ray flares with optical polarization angle (PA) swings and substantial polarization degree (PD) variations indicate the active role of the magnetic field during flares \citep[e.g.,][]{Marscher08,Abdo10,Blinov15}. Several models have been put forward to explain these phenomena \citep[e.g.,][]{Larionov13,Marscher14,ZHC15}, and a first-principle magnetohydrodynamics (MHD) based model is also under development \citep{ZHC16}. For the high energy emission, \cite{ZHC13} have shown that by combining the infrared/optical and the X-ray/$\gamma$-ray polarization signatures, it would be possible to distinguish the two models. Several X-ray and $\gamma$-ray polarimeters are currently proposed and/or under development \citep[e.g.,][]{Hunter14}. However, despite remarkable progress that has been made to improve these high-energy polarimeters, they commonly suffer from limited sensitivity. If the high-energy polarization signatures vary as rapidly as the low-energy (optical) polarization, it will be difficult for these polarimeters to measure, as they will integrate over episodes of vastly different PAs. This prompted us to investigate the time-dependent high-energy polarization signatures of lepto-hadronic blazar models in more detail. In this paper, we present a newly developed 3D multi-zone time-dependent hadronic model code, 3DHad. This new code is based on the one-zone time-dependent Fokker-Planck (FP) based lepto-hadronic code of \cite{Diltz15}, but generalized to 3D multi-zone. By coupling with the 3D polarization-dependent ray-tracing routines of the 3DPol code developed by \cite{ZHC14}, we will derive the time-dependent radiation and polarization signatures across the whole blazar SED, including all LTTEs. Hence, we can study the general phenomenology of the light curves and time-dependent polarization signatures. Hadronic models generally require very high jet powers and magnetic fields. Therefore, we will put physical constraints on the allowed parameter space by estimating the available jet power and magnetic field in the case of a Blandford-Znajek \citep{BZ77} powered jet. With the above consideration, we will predict detailed time-dependent polarization signatures from proton-synchrotron dominated hadronic models based on various jet conditions and flaring mechanisms. These results can be compared with multiwavelength light curves and future high-energy polarization measurements, putting stringent constraints on the blazar jet conditions in a hadronic model. We will describe our code setup and physical considerations in Section \ref{code}, sketch our model setup in Section \ref{model}, present case studies in Sections \ref{result1} and \ref{result2}, and discuss the results in Section \ref{discussion}.

} In this paper, we have presented the first 3D multi-zone time-dependent lepto-hadronic blazar code, 3DHad, describing a lepto-hadronic model in a parameter regime in which the high-energy emission is dominated by proton synchrotron radiation. By coupling with the 3DPol code, we are able to derive the time-dependent flux and polarization signatures of this lepto-hadronic blazar emission model, including all LTTEs. Our work thus makes the first attempt to study the time-dependent lepto-hadronic multi-wavelength polarization signatures of blazar emission. We have explicitly calculated the physical constraints for the hadronic model. Based on our estimates, if the Blandford-Znajek mechanism is responsible for powering the jet and providing the the magnetic field in the jet, the hadronic emission region cannot be very large due to the limited magnetic flux that the central black hole can provide. Therefore, the largest variability time scale in the observer's frame is unlikely to exceed a few days. Also, the high particle energy necessary for the lepto-hadronic scenario requires extreme jet powers. These constraints would also suggest that UHE extragalactic neutrinos are unlikely to be attributed to blazars, as photo-pion production is negligible. If the lepto-hadronic polarization signatures derived here are indeed detected in future observations, the required extremely efficient particle acceleration, strong magnetic field, and high jet power will seriously challenge our current understanding of AGN jet formation. We have demonstrated that the general time-dependent signatures of our proton-synchrotron dominated lepto-hadronic blazar model is dominated by the intrinsic time scale relations, namely, $t_{\rm ec} < t_{\rm lc} \lesssim t_{\rm pc}$. Through detailed parameter studies, we have identified the following time-dependent signatures of this model: \begin{enumerate} \item The time-dependent low-energy radiation signatures are generally symmetric in time, while the high-energy signatures are generally asymmetric; \item The high-energy flares generally peak later and last longer than the low-energy flares; \item An orphan flare in the high-energy component is possible; \item The polarization signatures at various wavelengths within the high-energy component are generally similar; \item In the quiescent state, if the low- and high-energy components are co-spatial, they share similar polarization degrees and angles. \item While the low-energy polarization signatures may vary rapidly during flares, high-energy polarization signatures appear generally stable. \item The time-dependent low-energy signatures and the high-energy polarization variations are generally synchronized with the disturbance propagation and the LTTEs. The high-energy flares, on the other hand, can last much longer due to the slow proton cooling. \end{enumerate} We suggest that these features can be tested with simultaneous multiwavelength observations, including future high-energy polarimetry. We notice that the polarization signatures possess a strong dependence on the magnetic field evolution. Although we have demonstrated in Section \ref{model} that our assumptions on the magnetic field evolution are reasonable, our test cases are most likely an over-simplification of any actual physical scenario. However, our code can be easily coupled with first principle simulations, such as MHD, to constrain the magnetic field evolution, so that our polarization signatures in both low- and high-energy components are physically self-consistent.

1607.06491

1607

1607.00879_arXiv.txt

The Narrow-line Seyfert 1 galaxy (NLS1) IRAS 13224$-$3809 is known to exhibit significant X-ray spectral variation, a sharp spectral drop at $\sim$ 7 keV, strong soft excess emission, and a hint of iron L-edge feature, which is very similar to the NLS1 1H 0707$-$495. We have proposed the ``Variable Double Partial Covering (VDPC) model'' to explain the energy spectra and spectral variability of 1H 0707$-$495 (Mizumoto, Ebisawa and Sameshima 2014, PASJ, 66, 122). In this model, the observed flux/spectral variations below 10 keV within a $\sim$day are primarily caused by change of the partial covering fraction of patchy clouds composed by double absorption layers in the line of sight. In this paper, we apply the VDPC model to IRAS 13224$-$3809. Consequently, we have found that the VDPC model can explain the observed spectral variations of IRAS 13224--3809 in the 0.5--10~keV band. In particular, we can explain the observed Root Mean Square (RMS) spectra (energy dependence of the fractional flux variation) in the entire 0.5 --10 keV band. In addition to the well-known significant drop in the iron K-band, we have found intriguing iron L-peaks in the RMS spectra when the iron L-edge is particularly deep. This feature, which is also found in 1H 0707--495, is naturally explained with the VDPC model, such that the RMS variations increase at the energies where optical depths of the partial absorbers are large. The absorbers have a larger optical depth at the iron L-edge than in the adjacent energy bands, thus a characteristic iron L-peak appears. On the other hand, just below the iron K-edge, the optical depth is the lowest and the RMS spectrum has a broad dip.

Among Active Galactic Nuclei (AGNs), Narrow-line Seyfert 1 galaxies (NLS1s) are characterized by their particular X-ray spectral and timing properties. A strong soft excess below $\sim$2 keV and remarkable X-ray variations are often observed, and high-energy spectral drops at $\sim$7 keV and seemingly broadened and skewed iron emission lines are found in several objects (e.g., \citealt{boller03}). These spectra are often explained by either the ``relativistic disk-line" model or the ``partial covering" model. According to the ``relativistic disk-line" model, their spectra may be interpreted by relativistically blurred inner-disk reflection around extreme Kerr black holes (e.g., \citealt{fabi04}). On the other hand, the ``partial covering" model may also explain their spectra as due to partial covering of the central X-ray source by intervening absorbers in the line of sight (e.g., \citealt{matsu90}; \citealt{ino03}; \citealt{millerl08}). Furthermore, \citet{nod11,nod13} suggests that the spectral continuum of AGNs may be more complex than previously considered. From the static spectral aspect alone, we cannot judge which model is more reasonable. To scrutinize the validity of these models in more detail, we need to explain not only the static spectral features but also their spectral variations. NLS1s are also characterized by significant X-ray time variation. In particular, their Root Mean Square (RMS) spectra (energy dependence of the fractional variation) tend to drop at the iron line energy band, which is most remarkably observed in the NLS1 MCG$-$6$-$30$-$15\footnote{Some authors argue that MCG$-$6$-$30$-$15 is a Seyfert 1 galaxy, but we treat it as NLS1 because it satisfies the properties of NLS1 \citep{Mchardy05}. } (\citealt{fabi02}; \citealt{matsumo03}). The ``relativistic disk-line" model explains the rapid spectral variations primarily by changes of the geometry in the very vicinity of the black hole, such as height of the illuminating source above the black hole (e.g., \citealt{mini04}). In this model, the disk-reflected photons are much less variable than the direct photons due to relativistic reverberation (e.g., \citealt{fabi03}), thus the characteristic RMS spectra are explained. In addition, \cite{fabi09} reported soft lags from the NLS1 1H 0707$-$495 and interpreted them as due to reverberation from the accretion disk, where the reflection component responds to variation in the X-ray corona with the corresponding time-lag. To the contrary, \cite{mil10} proposed that the soft X-ray lags of 1H 0707$-$495 can be accounted for by reverberation due to much more distant matter. Up to now, similar soft lags are detected in a number of NLS1s (e.g., \citealt{kara15}), but their origins are not fully understood. Meanwhile, \citet*{mizu14} (hereafter, Paper I) successfully explained the rapid variation of 1H 0707$-$495 by the ``variable double partial covering" (VDPC) model in the 0.5--10 keV. In this model, intrinsic X-ray luminosity and spectral shape of the central X-ray source below $\sim10$~keV are not significantly variable in timescales less than $\sim$day, and apparent X-ray variation is primarily caused by variation of the partial covering fraction by intervening absorbers composed of two different ionization layers. Spectral variations in $\sim2-10$ keV and the RMS spectra of 21 Seyfert galaxies including MCG$-$6$-$30$-$15 observed with Suzaku are also explained by the VDPC model successfully (\citealt{miya12,iso16}). We aim to examine whether the VDPC model can explain spectral variations of other NLS1s in wider energy ranges. In this paper, we apply the VDPC model to IRAS 13224$-$3809, which is characterized by the soft lag, significant time variation, a sharp spectral drop at $\sim$ 7 keV, strong soft excess emission, and a hint of an iron L-edge feature, being very similar to 1H 0707$-$495 (e.g., \citealt{gallo04}; \citealt{fabi13}).

We have studied spectral variations of NLS1 IRAS 13224$-$3809, using all the currently available XMM-Newton and Suzaku archival data. Following Paper I, we examined if the observed spectral variation is explained by the Variable Double Partial Covering (VDPC) model. Consequently, we have found that the VDPC model can successfully explain the averaged and intensity-sliced spectra of IRAS 13224$-$3809 in 0.5$-$10 keV within a $\sim$ day only changing the partial covering fraction. The model can explain the light-curves within a $\sim$day mostly by only change of the partial covering fraction, whereas some intrinsic variation above $\sim$ 3 keV is additionally recognized. We have successfully explained the observed RMS spectra in the entire 0.5$-$10 keV band with the VDPC model. In addition to the well-known significant drop in the iron K-band, we have found such intriguing broad iron L-peaks in the RMS spectra (as well as 1H 0707$-$495), that is particularly significant when the iron L absorption edge is deep in the energy spectra. These RMS spectral features can be explained by only change of the partial covering fraction, such that the RMS variation increases at the energies where the optical depth of the partial absorbers is large, and vice versa. The optical depth is minimum just below the iron K-edge and suddenly increases at the iron K-edge, thus the broad dip structure is produced. Around the iron L-energy band, the optical depth is the largest, thus the characteristic peak appears. \bigskip \bigskip This research has made use of public Suzaku data obtained through the Data ARchives and Transmission System (DARTS), provided by Institute of Space and Astronautical Science (ISAS) at Japan Aerospace Exploration Agency (JAXA). This work is also based on observations obtained with XMM-Newton, an ESA science mission with instruments and contributions directly funded by ESA Member States and the USA (NASA), and the XMM-Newton data obtained through the XMM-Newton Science Archive at ESA. For data reduction, we used software provided by the High Energy Astrophysics Science Archive Research Center at NASA/Goddard Space Flight Center. MM and KE are financially supported by the Japan Society for the Promotion of Science (JSPS) KAKENHI Grant Number 15J07567 and 16K05309, respectively. We acknowledge the referee, Dr.~J.~Reeves, for valuable comments.

1607.00879