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charlieoneill

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jsalt-astroph-dataset

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1012

1012.5770_arXiv.txt

We investigated the size distribution of high-inclination main-belt asteroids (MBAs) to explore asteroid collisional evolution under hypervelocity collisions of around 10 km s$^{-1}$. We performed a wide-field survey for high-inclination sub-km MBAs using the 8.2-m Subaru Telescope with the Subaru Prime Focus Camera (Suprime-Cam). Suprime-Cam archival data were also used. A total of 616 MBA candidates were detected in an area of 9.0 deg$^2$ with a limiting magnitude of 24.0 mag in the SDSS $r$ filter. Most of candidate diameters were estimated to be smaller than 1 km. We found a scarcity of sub-km MBAs with high inclination. Cumulative size distributions (CSDs) were constructed using Subaru data and published asteroid catalogs. The power-law indexes of the CSDs were 2.17 $\pm$ 0.02 for low-inclination ($< 15^{\circ}$) MBAs and 2.02 $\pm$ 0.03 for high-inclination ($> 15^{\circ}$) MBAs in the 0.7--50 km diameter range. The high-inclination MBAs had a shallower CSD. We also found that the CSD of S-like MBAs had a small slope with high inclination, whereas the slope did not vary with inclination in the C-like group. The most probable cause of the shallow CSD of the high-inclination S-like MBAs is the large power-law index in the diameter--impact strength curve in hypervelocity collisions. The collisional evolution of MBAs may have advanced with oligopolistic survival during the dynamical excitation phase in the final stage of planet formation.

The size distribution of main-belt asteroids (MBAs) traces their collisional history since the formation of primitive planetesimals. Previous survey observations for small solar system bodies have revealed the size distribution of MBAs down to sub-km size (Ivezi\'c et al. 2001, Yoshida \& Nakamura 2007, Gladman et al. 2009). Understanding the collisional processes of MBAs provides insight into the initial size distribution and collisional and dynamical evolution of planetesimals in the inner protoplanetary disk. Several models for collisional and orbital evolution of MBAs suggest a dynamical excitation event at the final stage of the planet formation \citep{P02}. This event produced several features seen in the current main belt, such as a high eccentricity/inclination asteroid population and the radial mixing of taxonomic classes. The dynamical excitation also ejected most asteroids out of the main belt, depleting more than 99\% of the primordial mass \citep{Bo05}. The relative velocities among asteroids were elevated during this phase. In the main belt, asteroid collisions mostly occurred at higher velocities than the present average velocity of $\sim$4 km s$^{-1}$ \citep{V98}. The relative velocity between a body in the remnant main belt population and one in the ejected population from the main belt zone exceeded 10 km s$^{-1}$ \citep{Bo05}. However, the fragment size distribution and ejecta velocities of such high-velocity collisions are still unclear; reproduction by laboratory experiments is difficult \citep{K10}. In our approach to this problem, we focus on MBAs with high inclination. The mean relative velocity of highly-inclined MBAs at collisions with other MBAs can be as large as $\sim$10 km s$^{-1}$ \citep{G06}. Comparing the size distributions between low-inclination (low-$i$) and high-inclination (high-$i$) MBAs could provide meaningful insight into disruptions at high velocity. For investigating high-$i$ MBAs, several published asteroid catalogs are available, but their detection limits are no smaller than kilometer size. We performed observations for sub-km high-$i$ MBAs to obtain their size distribution.

We performed wide-field observations of high-$i$ MBAs of sub-km diameter using the 8.2 m Subaru Telescope. The CSDs of low-$i$ and high-$i$ MBAs were constructed using the Subaru data and the ASTORB and SDSS MOC catalogs. A summary of our main conclusions follows. (1) A smaller fraction of small asteroids (0.7 km $<D<$ 1 km) appear in high inclination. The power-law index of the CSD for the 0.7--2 km diameter range is 1.79 $\pm$ 0.05 for low-$i$ MBAs ($i < 15^\circ$) and 1.62 $\pm$ 0.07 for high-$i$ ($i > 15^\circ$) MBAs. The sub-km asteroids at $i > 15^\circ$ have a lower index than the low-$i$ sample at the 95\% confidence level based in a statistical $t$-test. (2) The single power-law slope of the combined CSD for diameter in the 0.7--50 km range is 2.17 $\pm$ 0.02 for the low-$i$ MBAs and 2.02 $\pm$ 0.03 for the high-$i$ MBAs. The high-$i$ CSD is shallower than the low-$i$ CSD for the entire size range at the 99\% confidence level. This is not caused by the difference in the wavy pattern of the CSDs. (3) The CSD of S-like asteroids has a small slope in high inclination, whereas that of C-like asteroids shows little variation in slope with inclination. Through modeling we showed that the shallow CSD of the high-$i$ MBAs is not caused by the spatial distribution of the taxonomic groups, but by the shallow CSD of the S-like MBAs with high inclination. (4) The difference in slope of CSDs between the low-$i$ and high-$i$ MBAs is constant across $D \sim 10$ km. The shallow CSD of the high-$i$ MBAs is not the result of dynamical removal due to the Yarkovsky effect and the secular resonances. We suggest that the small $b$ of S-like MBAs with high inclination is due to a collisional effect. The possible explanation is that the $Q^{\ast}_{\rm D}$ curve has a large gravity-scaled regime $s_{\rm g}$ slope under hypervelocity collisions (around 10 km s$^{-1}$). Asteroid collisions often occurred with such high velocities in the dynamical excitation phase in the final stage of planet formation. We suppose that during this phase, MBAs experienced oligopolistic collisional evolution; small bodies were more easily disrupted relative to large bodies than at present. This indication claims that the current evolutionary models for MBAs should be modified. \bigskip

1012.5770

1012

1012.4540_arXiv.txt

The multifrequency radio continuum and 21cm \hi observations of five blue compact dwarf (BCD) galaxies, Mrk 104, Mrk 108, Mrk 1039, Mrk 1069 and I Zw 97 using the Giant Meterwave Radio Telescope (GMRT) are presented here. Radio continuum emission at 610 MHz and 325 MHz is detected from all the observed galaxies whereas only a few are detected at 240 MHz. In our sample, three galaxies (Mrk 104, Mrk 108 and Mrk 1039) are members of groups and two galaxies (Mrk 1069 and I Zw 97) are isolated galaxies. The radio emission from Mrk 104 and Mrk 108 is seen to encompass the entire optical galaxy whereas the radio emission from Mrk 1039, Mrk 1069, I Zw 97 is confined to massive \hii regions. This, we suggest, indicates that the star formation in the latter group of galaxies has recently been triggered and that the environment in which the galaxy is evolving plays a role. Star formation rates (SFR) calculated from 610 MHz emission is in the range $0.01-0.1\ M_\odot$~yr$^{-1}$; this is similar to the SFR obtained for individual star forming regions in BCDs. The integrated radio spectra of four galaxies are modelled over the frequency range where data is available. We find that two of the galaxies Mrk 1069 and Mrk 1039, show a turnover at low frequencies which is well fitted by free-free absorption whereas the other two galaxies, Mrk 104 and Mrk 108, show a power law at the lowest GMRT frequencies. The flatter spectrum, localized star formation and radio continuum in isolated galaxies lend support to stochastic self-propagating star formation (SSPSF). The \hi observations of four galaxies Mrk 104, Mrk 108, Mrk 1039 and Mrk 1069 show extended disks as large as $\sim1.1-6$ times the optical size. All the observed BCDs (except Mrk 104) show rotating disk with a half power width of $\sim50-124$ \kms. Solid body rotation is common in our sample. We note that the tidal dwarf (TD) origin is possible for two of the BCDs in our sample.

} \label{sec:intro_radio} Blue compact dwarf galaxies (BCDs) are star forming dwarf galaxies whose bluer colours are attributed to ongoing star formation. They are gas-rich, compact, low luminosity ($M_{B}= -17$ to $-14$) objects with low metal abundances ($\frac{1}{50}Z_{\odot}< \ Z <\frac{1}{2}Z_{\odot}$: \citealt{izo06}). They are not forming stars for the first time, as was predicted earlier (\citealt{ssb73}; \citealt{ss72}). All BCDs possess a faint low surface brightness component that is detected both in the optical and in the IR (\citealt{caon05} and references therein; \citealt{ramya09}), implying the presence of old red stars. By analyzing colour maps and surface brightness profiles of the low surface brightness (LSB) component, the ages and chemical abundances of the underlying host galaxies have been determined (\citealt{noeske00}; \citealt{pap02}; \citealt{ramya09}). The enrichment of the ISM in dwarf galaxies mainly occurs during these short starburst events \citep{leg00}. \cite{kun86} proposed that if the metals produced during a starburst are immediately mixed with the surrounding \hii regions, the metallicity will rise very quickly to values of the order of 1/50th of the solar value which explains why no galaxy with metallicity lower than I Zw 18 (1/50th $Z_\odot$), a nearby BCD, has ever been found \citep{leg00}. \\ A search for quiescent BCDs (QBCDs), carried out by \citet{alm08} indicates that after their bursting phase of a few 10 Myr to a few 100 Myr, BCDs enter the quiescent stage. BCDs spend about 30 times more time in the quiescent phase. However, QBCDs are found to be more metal rich than BCDs \citep{alm08} which is yet to be understood. It is noticed that, none of the dwarfs or low surface brightness (LSB) galaxies show a SFR equal to zero (\cite{leg00} and references therein) which implies even during quiescence star formation occurs at a very low rate. \cite{leg00} has concluded through his modelling that the observed oxygen and carbon abundances in I Zw 18 can be reproduced by a continuous SFR of $10^{-4}$ \msyr after 14 Gyr. However to reproduce the present colours, they had to include a bursting episode. All this suggests the existence of a weak but continuous regime of star formation in these galaxies. A study of extremely isolated BCDs by \cite{zitrin09} concludes that the galaxy colours are better explained by the combination of a continuous star formation process with a recent instantaneous star burst, than by a combination of several instantaneous bursts as suggested previously. \ A few BCDs also emit high ionization lines of \heii $\lambda4686$ \AA, \nev\ $\lambda\lambda3346,3426$ \AA, \fev $\lambda4227$ \AA, \fevii\ $\lambda\lambda5146,5177$ \AA, along with broad emission lines of H$\beta$, \oiii\ $\lambda\lambda4959,5007$ \AA \ and H$\alpha$ in very low metallicity and dense interstellar medium which are believed to be due to supernovae (SNe) events and/or stellar winds (\citealt{izo07}, and references therein). Whether these low-metallicity BCDs can host an active nucleus is another current area of research. Chandra X-ray observations of star bursting dwarf galaxies, such as NGC 1569 and NGC 3077 \citep{grimes05} seem to indicate that the material is blown out into the halo and consequently even removed from the galaxy leading to enrichment of the intergalactic medium. \ BCDs harbour appreciable amounts of dust \citep{tsm99}, confirmed from the far-IR (FIR) emission at 60 $\mu$m, $90\mu$m and $140\mu$m \citep{hirashita09}. The optical properties of dust are similar to the dust in the Milky Way \citep{hirashita09}. However the dust in BCDs appears to be warmer. Weak polycyclic aromatic hydrocarbons (PAH) emission in the bands at 6.2, 7.7, 11.2 and $12.8 \mu$m is detected in some BCDs. PAH emission is suppressed in most metal-poor BCDs, believed to be because of a metallicity threshold below which PAHs cease to form (\citealt{wu09}; \citealt{dwek05}). \cite{engel05} and \cite{engel08} found an anticorrelation between the dust temperature and metallicity implying warmer dust at lower metallicities of log(O/H)+12$\sim8$ and temperature continues to fall with further reductions in metallicity. The dependence on metallicity is found out to be $\sim Z^{-2.5}$ down to log(O/H)+12$\sim8$. The change in dust behaviour in terms of PAH emission, FIR colour temperatures and dust/gas mass ratio, all near metallicity log(O/H)+12$=8$ indicate that near this metallicity there is a general modification of the components of the interstellar dust that dominates the infrared emission \citep{engel08}. Radio observations which include the 21cm spectral line of \hi and radio continuum emission are useful in estimating the neutral gas content and kinematics, star formation rates and possible signatures of interactions. A sample of BCDs observed in \hi confirms that metal poor systems tend to be gas-rich low-luminosity galaxies \citep{hutchmeier07}. A range of spectral shapes at radio frequencies have been observed for BCDs (\citealt{hunt05}; \citealt{yin03}; \citealt{deeg93}; \citealt{kle91}). The observed radio continuum spectrum is attributed to star formation. The FIR-radio correlation of BCDs is similar to that of normal galaxies \citep{yun01}. The initial triggering mechanism, evolution of starburst and evolution of BCDs as a whole is not yet understood. Several mechanisms have been proposed, ranging from internal instabilities to external (especially tidal) triggers. If systems are isolated, star formation can be explained using the stochastic self propagating star formation mechanism (SSPSF), first proposed by \cite{gerola80}. Recent studies of large samples of star-forming dwarf galaxies \citep{noeske01} that look for faint companions support the hypothesis of interaction-induced star formation in BCDs. A lower limit for the fraction of star forming dwarf galaxies found with companions is $\sim30\%$ \citep{noeske01}. Thus, both the mechanisms are plausible and it is difficult to quantify the relative influence of the two mechanisms at a particular epoch. The five blue compact dwarf galaxies studied here (Mrk 104, Mrk 108, Mrk 1039, Mrk 1069 and I Zw 97) are selected from a larger sample chosen for an optical study (\citealt{ramya09}; \citealt{ram10}). In this paper, we present the 21cm \hi tracing the neutral atomic gas and radio continuum observations at 240, 325 and 610 MHz tracing the combined distribution of thermal and non-thermal radiation for the five BCDs. This is the first time that many of these galaxies have been detected at frequencies $<1$ GHz. Combined with the higher frequency observations from literature, where available, the radio spectra can be modelled. The distance to these galaxies is between 20 and 40 Mpc. Table \ref{tab:para} lists the general properties of these galaxies. Mrk 104 belongs to a loose group UZC-CG 94 consisting of 3 members, with the closest member, UGC 4906 an Sa galaxy separated from Mrk 104 in the sky plane by $\sim330$ kpc. The other member is PGC 26253. Mrk 104 has been classified as having a double nucleus in the process of merging \citep{mazarella93}. \cite{ramya09} resolve the two nuclei and note that both show \hii region like spectra thus ruling out the presence of an AGN in the centre of the galaxy. Mrk 104 is situated at a distance of 31.2 Mpc and at this distance $1\arcsec$ corresponds to $\sim151$ pc. Mrk 108 classified as an I0 in NED is one of the four members of the group Holmberg 124. NGC 2820 (of Hubble type SBc) is the closest neighbour. There is also a radio bridge connecting NGC 2820 and Mrk 108 to the third member of the group NGC 2814 \citep{nim05} clearly indicating a tidal interaction. Several other signatures of hydrodynamic processes are also observed in the group \citep{nim05}. This galaxy hosted a type IIp supernova, SN 1998bm indicating recent massive star production. Mrk 108 is situated at a distance of 22.2 Mpc and at this distance $1\arcsec$ corresponds to $\sim108$ pc. Mrk 1039 is a member of the group USGC S087 \citep{ramella02} and LGG 59 (Lyon group of galaxies: \citealt{garcia93}). DDO 023, DDO 020 and Mrk 1042 are the other members of the group. All the companion members are dwarf galaxies. Though located close to the Eridanus supergroup, it is not considered to be part of it \citep{brough06}. A type II supernova, SN 1985S is recorded in Mrk 1039. The galaxy is located at a distance of 28.8 Mpc (galactocentric distance taken from NED) and at this distance $1\arcsec$ corresponds to $\sim139$ pc. The fourth galaxy in our sample, Mrk 1069 also lies close to the Eridanus supergroup but is not considered to be member of the group. No group membership is assigned to Mrk 1069. This galaxy is located at a distance of 20.7 Mpc from us and at this distance $1\arcsec$ corresponds to a distance of $\sim100$ pc. I Zw 97 is an isolated galaxy with no neighbour within about $50\arcmin$ (525 kpc). \cite{thuan81} do not detect this galaxy in \hi and conclude that the atomic gas surface density is $< 2.7\times10^6\, M_\odot\,$Mpc$^{-2}$ and the upper limit on the \hi mass is $7.3\times10^8\, M_\odot$. Type II SN 2008bx was discovered recently in this galaxy, and \cite{atel09} reported the detection of radio continuum emission at 610 MHz from this supernova. The galaxy is at a distance of 36.1 Mpc. At this distance $1\arcsec$ corresponds to $\sim175$ pc.\\ This paper is structured as follows. Section \ref{sec:obs_rad} gives an account of the observations and data reduction. Section \ref{sec:notes_rad} gives a note on individual galaxies. Section \ref{sec:discus_rad} is a detailed discussion on our results and section \ref{sec:con_rad} summarizes the study. \begin{landscape} \begin{table*}[h!]\footnotesize \begin{center} \caption{\footnotesize The general parameters of the five galaxies collected from literature.} \begin{tabular}{llllll} \\ \hline\hline \textbf{Parameter} & \multicolumn{5}{c}{\textbf{Galaxy}} \\ \hline --- & \textbf{Mrk 104} & \textbf{Mrk 108} & \textbf{Mrk 1039} & \textbf{Mrk 1069} & \textbf{I Zw 97} \\ \hline \\ \textbf{Hubble type}$^e$ & Pec & I0 pec & Sc, edge-on, \hii & Sa$^*$ & --- \\ \textbf{Helio vel}$^e$ (\kms) & 2235 & 1534 & 2111 & 1562 $^d$ & 2518 \\ \textbf{Central vel}$^a$ (\kms) & 2203 & 1574 & 2098$^d$ & 1562$^d$ & 2530 \\ \textbf{Group} & UZC-CG 94$^b$ & Holm 124 & USGC S087$^c$ & --- & --- \\ \textbf{members} & UGC 4906, & NGC 2820, & DDO 023, & UGCA 052 & --- \\ --- & PGC 26253 & NGC 2814 & DDO 020 & --- & --- \\ --- & --- & --- & Mrk 1042 & --- & --- \\ \textbf{Single dish \hi mass}$^a$ (M$_\odot$) & $5.0\times10^8$ & $9.4\times10^9$ & $1.2\times10^9 \ ^d$ & $8.5\times10^8 \ ^d$ & {$<7.3\times10^8$}\ \\ \textbf{50\% line width}$^a$ (\kms) & 163 & 324 & 149 $^d$ & 106 $^d$ & --- \\ \textbf{Galactocentric Distance}$^{e,h}$ (Mpc) & 31.2 & 22.2 & 28.8 & 20.7 & 36.1 \\ \textbf{linear scale - $1''$}\ $^{e,h}$ (pc) & 151 & 108 & 139 & 100 & 175 \\ \textbf{m$_B$}$^f$ (mag) & $15.1\pm0.4$ & $15.5\pm0.3$ & 13.94$^g$ & 14.55$^g$ & $14.9\pm0.3$ \\ \textbf{M$_B$} (mag) & -17.46 & -16.47 & -18.46 & -17.32 & -16.76 \\ \textbf{L$_B$} ($10^9\times L_\odot$) & 1.5 & 0.6 & 3.8 & 1.3 & 0.79 \\ \textbf{L$_{FIR}$} ($10^9\times L_\odot$) & 0.54 & --- & 1.96 & 0.88 & 0.82 \\ \textbf{SN recorded} & --- & SN1998bm & SN1985S & --- & SN2008bx \\ \hline \end{tabular} \end{center} $^a$ - \cite{thuan81}, $^b$ - \cite{focardi02}, $^c$ - \cite{ramella02}, $^d$ - \cite{thuan99}, $^e$ - NASA Extragalactic Database, $^f$ - \cite{rc391}, $^g$ - \cite{doy05}, $^h$ - Assuming $H_0$ = 73 km~s$^{-1}$~Mpc$^{-1}$, * - Hyperleda. \label{tab:para} \end{table*} \end{landscape}

} \label{sec:con_rad} The radio continuum observations at frequencies 610 MHz, 240 MHz, 325 MHz, 1.28 GHz and 1.4 GHz with \hi observations and analyses of five blue compact dwarf galaxies (BCDs) using the Giant Meterwave Radio Telescope (GMRT) are presented here. The \hi observations of four BCDs namely Mrk 104, Mrk 108, Mrk 1039 and Mrk 1069 have revealed their large \hi disks about 1.1 -- 3 times the optical size of the galaxy. The typical \hi masses range between $2\times10^8-10^9\, M_\odot$. These values are typical of blue compact dwarf galaxies. We also detect a cloud close to the disk in Mrk 104 with no obvious optical counterpart and speculate that these might have influenced the current burst of star formation. Rotation is clearly noticed in the velocity maps of Mrk 108, Mrk 1039 and Mrk 1069. The \hi line profiles show multiple components. Mrk 104 has a \hi mass of $2\times10^8$, rotation velocity is 69 \kms \ and $M\textrm{(dyn)}/L_K\sim1$. Mrk 108 with total \hi mass is $1.6\times10^8\, M_\odot$, $M\textrm{(dyn)}/L_K\sim3$, rotation velocity $\sim52$ \kms \ and is situated very close to the galaxy NGC 2820 (a spiral galaxy). Mrk 104 and Mrk 108 are similar with the nearest group member being a large spiral. A tidal origin cannot be ruled out for these two galaxies. All the galaxies are detected in radio continuum at 325 MHz and 610 MHz and the emission at this frequency is dominated by non-thermal emission. Out of the five galaxies, only two galaxies namely Mrk 104 and Mrk 1069 are detected at 240 MHz. I Zw 97 is detected for the first time in the radio; neither the NVSS survey nor the VLA FIRST survey detect it. Thermal and non-thermal separation is attempted for the four galaxies namely, Mrk 104, Mrk 108, Mrk 1039 and Mrk 1069. While the observed spectrum of Mrk 1069 can be explained by synchrotron spectrum absorbed by thermal gas at the lower frequencies, the observed spectrum of Mrk 104 could be explained as being due to combined thermal and non-thermal emissions. The observed spectrum of Mrk 108 is only due to pure non-thermal emission and the continuum emission in Mrk 1039 is explained as a combination of thermal and non-thermal emission with the non-thermal emission being absorbed at the lower frequencies by the thermal gas mixed in the region. We estimate a thermal fraction at 1.4 GHz of $\sim80\%$ and $\sim45\%$ of the total emission for Mrk 104 and Mrk 1039 which is much higher than the $\sim10\%$ generally seen in normal spiral galaxies \citep{cond92}. The emission in Mrk 1039, I Zw 97 and Mrk 1069 is confined to the \hii region and we note that the epoch of star formation in Mrk 1039 is considered to be fairly young $\sim 4$ Myr \citep{huang99}. The SFR estimated from the observed emission at 610 MHz results in values in the range 0.01-0.1 \msyr. These are similar to the values of SFRs obtained for the individual star forming regions using H$\alpha$ fluxes \citep{ramya09} in case of Mrk 1039 and I Zw 97 whereas it matches with the SFR estimated from global \ha for Mrk 104. Emission detected at 610 MHz from Mrk 1039, Mrk 1069 and I Zw 97 is localized to the \hii regions seen in the H$\alpha$ images, which indicates that the SFRs obtained from the radio measurement represents the SFRs of a few individual \hii regions in these galaxies. On the other hand, the 610 MHz emission from Mrk 104 is seen to encompass the entire galaxy. Combining our results with those of \cite{deeg93} --- the only other study of BCDs in low frequency radio continuum --- we find that on the average, the galaxies which are classified as isolated galaxies tend to show a somewhat flatter observed spectrum as compared to galaxies which are classified as being in groups. This, we interpret as due to a larger fraction of thermal emission mixed with the non-thermal in the former case, as one expects for young localized starbursts. However a flatter injection spectrum of a young localized starburst region could also explain this. We further note that the starburst in a galaxy in tenuous environment is likely to be more localized. The flatter spectrum seen in isolated galaxies with localized star formation in the form of compact star forming regions and localized radio continuum emission, all suggest stochastic self-propagating star formation.

1012.4540

1012

1012.0891_arXiv.txt

We present $J$-band long-slit spectroscopic observation of NGC 1068 classified as a Seyfert 2 galaxy. $J$-band observations with OAO/ISLE provide clear detection of spatially extended [Fe II]1.257$\mu$m and [P II]1.188$\mu$m lines. We found that [Fe II]1.257$\mu$m/[P II]1.188$\mu$m increases with distance from a central continuum peak. Observed line ratios around the nucleus (continuum peak) are consistent with a typical value expected from photoionization models, while the ratios at 3$\arcsec - 4\arcsec$ ($210-280$ pc) east and west of the nucleus are slightly higher than this. In the off nucleus region of NGC 1068 we also found a possible association between [Fe II]1.257$\mu$m/[P II]1.188$\mu$m and the radio continuum. This suggests a mild contribution of shock ionization induced by a radio jet outside nucleus while photoionization by the central energy source is dominant near the nucleus.

\label{introduction} The narrow line regions (NLRs), which extend to several hundred or kilo parsec scale around the galaxy center, are the exclusive structure of active galactic nuclei (AGN) where the spatially resolved observations are possible. Therefore NLRs are often investigated as an important tool to study the ionization state of the interstellar medium (ISM) and/or chemical evolution in galactic scale (e.g., \cite{2006A&A...447..863N}). Although it is widely accepted that the NLR is photoionized by ionizing photons radiated from a central engine, the possibility of shock ionization induced by a jet in off nucleus regions cannot be excluded since ionization photons decrease with distance from a nucleus (e.g., \cite{2007ApJ...666..794F}). Thus, how much the shock contributes to the ionization of NLRs is very important to our understanding of AGN structure and in examining the utility of NLRs as a tool to investigate galactic-scale phenomena. Furthermore, the shock ionization of NLRs is getting a lot more attention. Recent dramatic progress of theoretical simulations and observational studies of galaxy formation and evolution allow a quantitative comparison between both sides. In this context, a serious problem has arisen, i.e., theoretical simulations predict too many massive galaxies due to long-duration star formations in contrast to early-time quenching of star formation in observed massive galaxies (e.g., \cite{2006MNRAS.365...11C,2006MNRAS.370..645B}). This problem cannot be solved even if a negative feedback effect on star formation activity caused by supernovae is involved and so AGN feedback effect is considered as a potential candidate of a solution: a massive galaxy likely has a supermassive black hole at its galaxy center, and inflow of ISM to a supermassive black hole invokes its AGN activity which releases vast gravitational potential energy to ISM resulting in suppression of star formation activity (e.g., \cite{2005ApJ...635L..13S,2007MNRAS.380..877S}). However, how the AGN activity transmits its energy to ISM remains a mystery. One possible physical mechanism of the AGN feedback is shock ionization of ISM, i.e., the AGN activity inputs its energy to ISM through a shock heating induced by a jet. In previous studies of NLRs, line-ratio diagnostics to distinguish between shock ionization and photoionization have been examined. Optical diagnostics, however, can hardly discriminate between the two mechanisms, because optical NLR spectra predicted by photoionization and shock ionization models are very similar to each other \citep{1995ApJ...455..468D,1996ApJS..102..161D,2008ApJS..178...20A}. The near infrared line ratio of [Fe II]1.257$\mu$m/[P II]1.188$\mu$m is one of the most powerful indicators to discriminate photoionization and shock ionization. Both lines have similar critical densities and excitation temperatures, i.e., this line ratio is roughly proportional to the ratio of gas-phase abundance of iron and phosphorous. In contrast, iron is a well known refractory species and is strongly depleted in dust grains, whereas phosphorous is a non-refractory species. Photoionization alone (including H II regions and NLRs excited by ionizing photons from young stars and AGN central sources, respectively) is relatively incapable of destroying the tough iron based grains, while these are easily sputtered by shocks. The [Fe II]1.257$\mu$m/[P II]1.188$\mu$m ratio, therefore, is high ($\gtrsim 20$) in fast shock-excited regions and low ($\lesssim 2$) in normal photoionized regions \citep{2001A&A...369L...5O}. The actual ionization state of NLRs would be determined by the combination of photoionization and shock ionization, and so the observed line ratios are expected to change with the locations in NLRs ranging from [Fe II]1.257$\mu$m/[P II]1.188$\mu$m $\sim 2$ to $\sim 20$. NGC 1068 is one of the nearest AGNs, which has a compact radio jet around the nucleus and spatially extended radio lobe \citep{1983ApJ...275....8W}. This radio structure well coincides with a morphology of NLR \citep{1997ApJ...487..560C}. Therefore, NGC 1068 is an ideal object with which to research the spatial distribution of [Fe II]1.257$\mu$m/[P II]1.188$\mu$m. In this paper, we adopt 14.4 Mpc as a distance to NGC 1068 and 1$\arcsec$ corresponds to 70 pc.

The line ratio [Fe II]1.257$\mu$m/[P II]1.188$\mu$m in the near-infrared wavelength range is a useful tool with which to examine the dust destruction by shocks. We investigated spatial distribution of this ratio in NLR of nearby Seyfert galaxy NGC 1068 with OAO/ISLE. [Fe II]1.257$\mu$m/[P II]1.188$\mu$m near the nucleus is close to unity consistent with a previous observation and with a ratio in a normal photoionized region. This indicates that photoionization by ionizing photons radiating from a central engine is dominant near the nucleus. We found that the ratio increases with the distance from the nucleus, and is slightly higher at 3$\arcsec$ $-$ 4$\arcsec$ east and west of the nucleus than ratios typical of a photoionized region. We also found a possible spatial association between [Fe II]1.257$\mu$m/[P II]1.188$\mu$m and radio continuum around $\sim 560$ pc from the nucleus. These findings suggest a higher contribution of shock ionization induced by a radio jet at off nucleus. Except for NGC 1068, recently the spatial correlation between [Fe II]1.257$\mu$m/[P II]1.188$\mu$m and radio continuum over the several hundred parsec scale has been reported for NGC 4151 and Mrk 1066. Applying this kind of research to a number of other AGNs is the clue to revealing ongoing AGN feedback phenomena. \\ We would like to thank Nozomu Kawakatu for his meaningful comments on interpretation of observed data. This work was supported by the Publications Committee of the National Astronomical Observatory of Japan (NAOJ) and the Grant-in-Aid for the Global COE Program \lq\lq The Next Generation of Physics, Spun from Universality and Emergence\rq\rq from the Ministry of Education, Culture, Sports, Science and Technology (MEXT) of Japan. T.N. acknowledges financial supports through the Research Promotion Award of Ehime University and the Kurata Memorial Hitachi Science and Technology Foundation. K.M. acknowledges financial support from the Japan Society for the Promotion of Science (JSPS) through the JSPS Research Fellowships for Young Scientists. \onecolumn \begin{figure} \begin{center} \FigureFile(80mm,80mm){figure-1.eps} \end{center} \caption{ $J$-band image of NGC 1068 obtained with OAO/ISLE in our observation. The long-slit position (P.A. = 90$^{\circ}$) is shown by two solid lines. }\label{figure1} \end{figure} \begin{figure} \begin{center} \FigureFile(160mm,80mm){figure-2.eps} \end{center} \caption{ 2-D spectra in $J$ band extracted from central $\pm$15$\arcsec$ region (a) and continuum-subtracted spectrum (b). }\label{figure2} \end{figure} \onecolumn \begin{figure} \begin{center} \FigureFile(80mm,80mm){figure-3.eps} \end{center} \caption{ [Fe II]1.257$\mu$m/[P II]1.188$\mu$m line ratio (top) and VLA 4.86 GHz flux density (bottom) as a function of distance from a continuum peak. Arrows in the top figure are lower limits calculated from 3 $\sigma$ noise level around undetected [P II]1.188$\mu$m. }\label{figure3} \end{figure} \begin{table} \caption{Relative line fluxes normalized by [P II]1.188$\mu$m}\label{table} \begin{center} \begin{tabular}{lclclc|c|} \hline Line ID&East 3$\arcsec.0$&Central 2$\arcsec$.0&West 3$\arcsec$.0\\ \hline ${\rm [P\ II]}$1.188$\mu$m&1.0$\pm$0.29&1.0$\pm$0.02&1.0$\pm$0.08\\ ${\rm [S\ IX]}$1.252$\mu$m&1.05$\pm$0.02&1.01$\pm$0.02&0.73$\pm$0.07\\ ${\rm [Fe\ II]}$1.257$\mu$m&1.79$\pm$0.02&1.33$\pm$0.05&1.63$\pm$0.07\\ Pa$\beta$&2.35$\pm$0.04&2.89$\pm$0.05&2.55$\pm$0.08\\ \hline \end{tabular}\\ Spectra were extracted from central 2$\arcsec$.0 region and east and west neighbor 3$\arcsec$.0 regions. \end{center} \end{table} \bigskip \newpage

1012.0891

1012

1012.2770_arXiv.txt

We have investigated the relaxation of a hydrostatic hot plasma column containing toroidal magnetic field by the Current-Driven (CD) kink instability as a model of pulsar wind nebulae. In our simulations the CD kink instability is excited by a small initial velocity perturbation and develops turbulent structure inside the hot plasma column. We demonstrate that, as envisioned by Begelman, the hoop stress declines and the initial gas pressure excess near the axis decreases. The magnetization parameter $\sigma$, the ratio of the Poynting to the kinetic energy flux, declines from an initial value of $0.3$ to about $0.01$ when the CD kink instability saturates. Our simulations demonstrate that axisymmetric models strongly overestimate the elongation of the pulsar wind nebulae. Therefore, the previous requirement for an extremely low pulsar wind magnetization can be abandoned. The observed structure of the pulsar wind nebulae do not contradict the natural assumption that the magnetic energy flux still remains a good fraction of the total energy flux after dissipation of alternating fields.

The pulsar wind nebulae (PWNe) may be considered as a relativistically hot bubble continuously pumped by an electron-positron plasma and magnetic field emanating from the pulsar. Pulsars lose their rotational energy predominantly by generating a highly magnetized, ultrarelativistic wind. The wind presumably terminates at a strong reverse shock and the shocked plasma inflates a bubble within the external medium. The synchrotron and inverse Compton radiation from the shocked plasma is observed from the radio to the gamma-ray band (see the review by Gaensler \& Slane 2006). Close to the pulsar the energy is carried mostly by electro-magnetic fields as Poynting flux; however, the common belief is that at the termination shock the wind must already be very weakly magnetized. Simple spherical models of PWNe suggest that the magnetization parameter $\sigma$, the ratio of the Poynting to the kinetic energy flux, needs to be as small as 0.001-0.01 just upstream of the termination shock (Rees \& Gunn 1974; Kennel \& Coroniti 1984a,b; Emmering \& Chevalier 1987). The reason for the required low value of $\sigma$ at the termination shock is that conservation of the magnetic flux in the effectively incompressible subsonic flow downstream of the termination shock implies rapid increase in the magnetic field strength with radius and the field within the nebula could exceed the equipartition value if the magnetization at the termination shock is not extremely small. Extensive axisymmetric MHD simulations of the flow produced by the pulsar wind within a plerionic nebula (Komissarov \& Lyubarsky 2003, 2004; Del Zanna et al.\ 2004, 2006; Volpi et al.\ 2008; Camus et al.\ 2009) show that one can account for the morphology of PWNe, including the remarkable jet-torus structure, with $\sigma\approx 0.01$. If the magnetization were larger, the nebula would be elongated by the magnetic pinch effect beyond observational limits. Such a low value of $\sigma$ is puzzling because it is not easy to invent a realistic energy conversion mechanism to reduce $\sigma$ to the required level. This problem, often referred to as the ``$\sigma$ problem'' is widely discussed in the literature (see recent reviews by Arons 2007; Kirk et al.\ 2009). One has to stress that all the available observation limits on $\sigma$ are obtained from the analysis of the plasma flow and radiation beyond the termination shock, where the upstream $\sigma$ is calculated from the ideal MHD jump conditions as if the Poyntning flux is transferred by large scale magnetic fields. However in the pulsar wind, most of the energy is transferred by waves, which an obliquely rotating magnetosphere excites near the light cylinder. These waves cannot propagate within the nebula because the wavelength (on the order of the light cylinder radius) is less than the particle Larmor radii. The above mentioned observational limits on $\sigma$ refer only to the mean magnetic field remaining after the oscillating part is erased. In the equatorial belt of the wind, the sign of the magnetic field alternates with the pulsar period, forming stripes of opposite magnetic polarity (Michel 1971; Bogovalov 1999); such a structure is called a striped wind. In the striped wind, Poynting flux can be converted into particle energy flux when the oppositely directed magnetic fields annihilate. Observations of X-ray tori around pulsars (Gaensler \& Slane 2006) as well as theoretical modeling of the pulsar wind (Bogovalov 1999; Spitkovsky 2006) suggest that it is in the equatorial belt where most of the wind energy is transported. Therefore, in the equatorial belt magnetic dissipation of the striped wind is the main energy conversion mechanism in pulsars. It has been shown that due to relativistic time dilation, complete dissipation could occur only on a scale comparable to or larger than the radius of the termination shock (Lyubarsky \& Kirk 2001; Kirk \& Skjaeraasen 2003, Zenitani \& Hoshino 2007). However, the alternating fields still annihilate at the termination shock (Petri \& Lyubarsky 2007). At higher latitudes, where the magnetic field does not change sign, fast magnetosonic waves may transport a significant amount of energy. These waves can decay relatively easily (Lyubarsky 2003) but can release only a fraction of the Poynting flux into the plasma, because at these latitudes most of the energy is carried by the mean magnetic field. The fraction of the energy transferred by the mean field can be found only from 3D numerical simulations of the pulsar magnetosphere. Even though this fraction is still not known, this fraction is less than 1/2 because the angular distribution of the Poynting flux in the pulsar wind is a maximum at the rotational equator, where the mean field is zero. This suggests that $\sigma$ becomes less than unity after the waves decay. Therefore, at a quantitative level the $\sigma$ problem is partially solved if Poynting flux is converted into plasma energy via dissipation of the oscillating part of the field. However, the residual $\sigma$ still cannot be as small as required by axisymmetric models. Therefore the question still remains as to how the mean field $\sigma$, which is only somewhat less than unity, could become extremely small. Since no mechanism had been found for the extraction of energy from a large scale, axisymmetric magnetic field, Begelman (1998) suggested that the problem can be alleviated if a current-driven (CD) kink instability destroys the concentric field structure in the nebula. In the axisymmetric case, magnetic loops in the expanding flow are forced to expand and perform work against the magnetic tension. The CD kink instability allows the loops to come apart and one expects that in 3D, the mean field strength is not amplified much by expansion of the flow and the hoop stress would not necessarily pinch the flow as much as would otherwise be supposed. In this case $\sigma$ just upstream of the termination shock might not need to be so unreasonably small as was found in axisymmetric models. This idea can be checked only by 3D simulations of plasma flow within the nebula. As a first step we simulate the 3D evolution of the simple cylindrical model of PWNe developed by Begelman \& Li (1992). This model describes a quasi-static cylindrical configuration with a purely toroidal magnetic field. The plasma within the cylinder is relativistically hot and the hoop stress is balanced by the thermal pressure. The cylinder is confined on the outside by a nonmagnetized gas. The linear analysis shows (Begelman 1998) that such a configuration is unstable with respect to the CD kink instability. Here we simulate the nonlinear evolution of this system in order to see whether stabilizing boundary conditions suppress development of the instability inside the plasma volume via 3D relativistic magnetohydrodynamic (RMHD) simulations of the CD kink instability in a relativistically hot plasma column containing a toroidal magnetic field. This paper is organized as follows: We describe the numerical method and setup used for our simulations in \S 2, present our results in \S 3, and discuss the astrophysical implications in \S 4.

We have investigated the development of the CD kink instability of a hydrostatic hot plasma column containing a toroidal magnetic field as a model for pulsar wind nebulae. The CD kink instability is excited by a small initial velocity perturbation and turbulent structure develops inside the hot plasma column. At the end of nonlinear evolution of the CD kink instability, the hot plasma column relaxes with a slow radial expansion. The magnetization $\sigma$ decreases from an initial value of $0.3$ to a final value of $0.01$. We find that for different initial perturbation profiles the time scale is a bit different but the physical behavior is the same. Therefore the relaxation of a hot plasma column is independent of the initial perturbation profile. Our simulations confirm the scenario envisaged by Begelman (1998). Toroidal magnetic loops come apart, the hoop stress declines and the pressure difference across the nebula is washed out. In our simulations, the ratio of the gas pressure on the axis to the total (magnetic+gas) pressure at the plasma column boundary decreases from 2.5 to 1.5 during the linear phase, while magnetic dissipation is still small. In the nonlinear phase, the magnetic field dissipates and the gas pressure excess near the axis disappears. For this reason, elongation of a pulsar wind nebula cannot be correctly estimated by axisymmetrical models, because axisymmetric models retain a concentric toroidal magnetic field geometry. Radiation from the Crab nebula is highly polarized along the axis of the nebula, which is indicative of a toroidal magnetic field. We see that even though the instability eventually destroys the toroidal structure, the magnetic field becomes completely chaotic only at the end of the nonlinear stage of development. Therefore, the toroidal magnetic field should dominate in the central parts of the nebula that are filled by recently injected plasma. Our simple model does not allow us to determine the value of $\sigma$ in the pulsar wind; for this purpose one has to perform fully 3D simulations of the expanded nebula taking into account the continuous injection of plasma. Here we have demonstrated that 3D effects are crucially important to a determination of the structure of pulsar wind nebulae and that previous dynamical arguments concluding that $\sigma$ must be extraordinarily small can be abandoned.

1012.2770

1012

1012.4010_arXiv.txt

We present the results of an optical spectroscopic study of 12 {\it GALEX}--discovered star-forming \lya~emitting galaxies (LAEs) at z $\sim$ 0.3. We measure the emission line fluxes from these galaxies by fitting their observed spectra to stellar population models in order to correct for underlying stellar absorption. We revisit earlier stellar population model fitting results, finding that excluding now-known AGNs lowers the typical stellar population age and stellar mass of this sample to $\sim$ 300 Myr and $\sim$ 4 $\times$ 10$^{9}$ $\mathcal{M}$\sol, respectively. We calculate their dust extinction using the Balmer decrement, and find a typical visual attenuation of A$_\mathrm{V}$ $\sim$ 0.3 mag, similar to that seen in many high-redshift LAEs. Comparing the ratio of \lya/H$\alpha$ and the \lya\ equivalent widths to the measured dust extinction, we find that the ISMs in these objects appear to be neither enhancing nor seriously attenuating the \lya\ equivalent widths, as would be the case in a quasi-clumpy ISM. Lastly, we perform a detailed analysis of the gas--phase metallicities of these galaxies, and we find that most galaxies in our sample have $Z$ $\lesssim$ 0.4 $Z$\sol. We find that at a fixed stellar mass, these low-redshift LAE analogs are offset by $\sim$ 0.6 dex lower in metallicity from the general galaxy population at similar redshifts based on the local mass-metallicity relationship. This implies that galaxies with \lya\ in emission may be systematically more metal poor than star-forming galaxies at the same stellar mass and redshift, similar to preliminary results at $z \sim$ 2.

Galaxies selected on the basis of a bright \lya~emission line were originally thought to be indicative of primordial galaxies undergoing their first burst of star--formation \citep{partridge67}, though they have recently been shown to be a complicated group of objects. Studies utilizing the technique of spectral energy distribution (SED) fitting have shown that typical narrowband--selected \lya~emitting galaxies (LAEs) appear to be predominantly young and low--mass, with ages $<$ 100 Myr and masses $\lesssim$ a few $\times$ 10$^{9}$ $\mathcal{M}$\sol\ \citep[e.g.,][]{gawiser06a, finkelstein07, pirzkal07, lai07, nilsson07a}. However, in many cases these galaxies also appear to contain some dust, thus they are likely not primordial in nature \citep[e.g.,][]{pirzkal07, lai08, finkelstein08, finkelstein09a, pentericci09}. In addition, a small fraction of LAEs appear to be more evolved with ages of $\sim$ 0.5 Gyr and masses of $\sim$ 10$^{10}$ $\mathcal{M}$\sol, suggesting that there may be multiple populations of LAEs, or perhaps a tail in the distribution of LAE properties toward more evolved objects \citep[e.g.,][]{finkelstein09a, lai08, pentericci09}. These more evolved LAEs provide a link to the characteristically more evolved Lyman break galaxies \citep[LBGs; e.g.,][]{kornei10}. While the stellar masses of galaxies can be reasonably constrained from SED fitting, the remaining properties suffer from degeneracies, and can be poorly constrained in the absence of rest--frame optical detections, especially at high--redshift. Specifically, the derived ages, dust extinctions and metallicities are only weakly constrained, as they all result in a reddening of the integrated light from a given galaxy. While photometry spanning the 4000 \AA\ break can improve the fidelity of age estimates, the extinction can typically only be roughly constrained, and the metallicity not constrained at all. In order to obtain more robust estimates of the physical make--up of LAEs, direct measurements of the extinction and metallicity are necessary. At low redshift, this is typically done using measurements of the flux of nebular emission lines in the rest--frame optical, e.g., using Balmer line ratios to measure the dust extinction, and ratios of metal lines to measure the gas--phase metallicities \citep[e.g.][]{kobulnicky99, pettini04}. These analyses are not currently possible at z $\gtrsim$ 3, as these diagnostic lines are shifted into the mid--infrared. From 2 $<$ z $<$ 3, these measurements are possible using near--infrared spectroscopy on 8--10m class telescopes \citep[e.g.,][]{erb06b, finkelstein10d}. However, samples of LAEs at these redshifts are only now being compiled, and the required integration times are long. Presently, we can use low--redshift analogs to try to understand LAEs at high--redshift. Locally, the \lya\ properties of star-forming galaxies have been studied in great detail. It was found that \lya\ emission can correlate with bright UV (and H$\alpha$) emission in some regions of a galaxy, but not in others. It was also shown that \lya\ can escape from galaxies even with a significant dust content, especially if outflows are present in the interstellar medium \citep[ISM; e.g.][]{kunth98, hayes07, atek08, ostlin09}. These results indicate that the mechanisms which regulate \lya\ escape are complicated, which is an important detail to consider when interpreting the results of high-redshift LAEs. A little further out, \citet{deharveng08} published the discovery of $\sim$ 100 LAEs at $z \sim$ 0.3 using the space--based {\it Galaxy Evolution Explorer} ({\it GALEX}) telescope. A study of their luminosity function found that LAEs at low redshift are more rare and less luminous in \lya\ than at $z >$ 3 \citep{deharveng08, cowie10}. In a followup study, \citet{finkelstein09c} studied the stellar populations of a subsample of 30 of these LAEs, finding that they appear older and more massive than typical high-redshift LAEs. A higher fraction of low-redshift LAEs appear to host active galactic nuclei (AGN), from $\sim$ 15 -- 40\% \citep{finkelstein09e, scarlata09, cowie10}. \citet{cowie10} compare the metallicities from the ratio of [N\,{\sc ii}]/H$\alpha$ emission from $z \sim$ 0.3 LAEs to those of $z \sim$ 0.3 continuum-selected galaxies, finding that while the distributions overlap, the LAEs extend to lower metallicties. Lastly, \citet{atek09} used H$\alpha$ and H$\beta$ observations from a subset of these galaxies to study the \lya\ escape fraction, finding clear evidence for a decrased escape fraction with increased extinction, although some galaxies do exhibit a \lya\ escape fraction greater than the continuum, which could imply a clumpy ISM geometry \citep[e.g.][]{neufeld91, finkelstein09a}. In this Paper we present the spectra of our sample of 12 {\it GALEX} LAEs in the Extended Groth Strip (EGS) which are dominated by star-formation activity, first analyzed in \citet{finkelstein09c} via SED-fitting, and later shown to be devoid of AGN activity in \citet{finkelstein09e}. In \S 2 of this Paper, we present the full spectra of every object in our sample, along with their measured line fluxes. In \S 3, we revisit the typical ages and masses of this sample, comparing to what we earlier derived in \citet{finkelstein09c} when knowledge of AGN activity was unknown. We discuss the dust properties and insights into the ISM geometries of our sample in \S 4, and in \S 5, we present metallicity measurements with three separate metallicity indicators, as well as examine the mass-metallicity relation, and the implications it has on LAEs near and far. In \S 6 we present our conclusions. Where applicable, we assume H$_\mathrm{o}$ = 70 km s$^{-1}$ Mpc$^{-1}$, $\Omega_{m}$ = 0.3 and $\Omega_{\Lambda}$ = 0.7.

1012.4010

1012

1012.3406_arXiv.txt

We present the photometric catalogs for the star-forming cluster NGC~602 in the wing of the Small Magellanic Cloud covering a range of wavelengths from optical ({\it HST/ACS} F555W, F814W and SMARTS/ANDICAM {\it V}, {\it I}) to infrared (\spit\ /IRAC 3.6, 4.5, 5.8, and 8~$\mu$m and MIPS 24~$\mu$m). Combining this with IRSF (InfraRed Survey Facility) near-infrared photometry ({\it J}, {\it H}, $K_s$), we compare the young main sequence (MS) and pre-main sequence (PMS) populations prominent in the optical with the current young stellar object (YSO) populations revealed by the infrared (IR). We analyze the MS and PMS population with isochrones in color-magnitude diagrams to derive ages and masses. The optical data reveal $\sim 565$ PMS candidates, low mass Stage~{\sc iii} YSOs. We characterize $\sim 40$ YSOs by fitting their spectral energy distributions (SEDs) to a grid of models \citep{robitaille07} to derive luminosities, masses and evolutionary phase (Stage~{\sc i}-{\sc iii}). The higher resolution {\it HST} images reveal that $\sim70\%$ of the YSO candidates are either multiples or protoclusters. For YSOs and PMS sources found in common, we find a consistency in the masses derived. We use the YSO mass function to derive a present-day star-formation rate of $\sim 0.2-1.0 ~\Msun$ yr$^{-1}$ kpc$^{-2}$, similar to the rate derived from the optical star formation history suggesting a constant star formation rate for this region. We demonstrate a progression of star formation from the optical star cluster center to the edge of the star forming dust cloud. We derive lifetimes of a few $10^5$ years for the YSO Stages~{\sc i} and {\sc ii}. Please note that color images are compressed for space. Please contact the first author for superior, high resolution versions.

\label{intro} The Small Magellanic Cloud (SMC) is a valuable astrophysical laboratory for understanding the processes of star formation in a galaxy that is extremely different from the Milky Way. In particular, it has a subsolar chemical abundance of Z$\sim 0.004$ \citep{rolleston99,lee05} and a low dust-to-gas ratio of $\sim$~1/30 Milky Way in the diffuse interstellar medium \citep{stanimirovic00} and $\sim1/6$ in star-forming regions \citep{bot07}. The SMC's lack of organized rotation means that star formation is predominantly driven by a combination of tidally-induced cloud-cloud interactions \citep{zaritsky00,zaritsky04} and shell formation \citep{hatzidimitriou05}. With its close proximity \citep[60.6~kpc;][]{ hilditch05}, the SMC is uniquely suited to detailed investigation of the stellar content, down to the sub-solar mass regime, in the youngest and most compact star clusters. The small, young star cluster NGC~602, associated with the highly structured H~{\sc ii} region N90 \citep{henize56}, is one of the most interesting star forming regions in the SMC. Located at the boundary between the SMC wing and the Magellanic Bridge, at the intersection of three H{\sc i} shells, it is probably the result of the interaction of two expanding shells of H{\sc i} that occurred approximately $\sim 7\, {\rm Myr}$ ago \citep{nigra08}. \citet{nigra08} also show that the N90 region is quiescent, with negligible H$\alpha$ shell expansion velocities. Propagation of star formation is most likely driven by radiation with stars forming along the edges of the photodissociation region. Star formation started approximately 4~Myr ago, with the formation of the central cluster, and gradually propagated towards the outskirts, where star formation continues \citep{carlson07}. \citet{gouliermis07} provide a list of 22 candidate Young Stellar Objects (YSOs) in the outskirts of N90. The optical cluster's Mass Function (MF) is consistent with a standard Salpeter~\citep{salpeter55} Initial Mass Function \citep{schmalzl08}. \citet{cignoni09} reconstruct a complete star formation history for the optical population and find that the pre-main sequence (PMS) is not more than $\sim$5~Myr, and the star formation rate (SFR) has reached a maximum in the last 2.5~Myr. We employ a new panchromatic approach to extend the star formation history analysis to the present day. Population identification via Color-Magnitude Diagram (CMD) analysis informs us of cluster-scale star formation, with optical revealing a bright main sequence (MS) and a young PMS and infrared (IR) CMD highlighting the youngest embedded sources. Using spectral energy distribution (SED) analysis, for which we combine Hubble Space Telescope ({\it HST}) optical to probe central stellar sources and Spitzer Space Telescope (\spit) IR to probe circumstellar disks and envelopes, we characterize star formation in the NGC~602 region (Figure~\ref{8col}) on the scale of single sources and protoclusters. The structure of this paper is as follows. We describe observations and data reduction in Section \ref{obs}. In Section~\ref{cmdanal}, we discuss what can be gleaned from the IR and optical data separately via CMD analysis. We describe the application of the SED fitter in Section~\ref{FIT} and present the results of our SED analysis of 77 sources, combining optical to mid-IR data, for YSOs (\ref{PAHs}, \ref{yseds0}, \ref{yseds1}, and \ref{yseds+}) and other types of IR sources (\ref{nseds}). We address the YSO Mass Function and Star Formation Rate (SFR) in Section \ref{mf} and the spatial and temporal distribution of young sources in Section~\ref{SpaceTime}.

We present a multi-wavelength analysis of photometry and imaging of the NGC~602 active star-forming region in the SMC, covering {\it HST} 0.55 through \spit\ 24~$\mu$m. From these data, we define stellar and proto-stellar populations and their spatial distribution. We estimate the present-day star formation rate and derive time scales for the formation of Stage~{\sc i} and Stage~{\sc ii} YSOs. We provide full mutli-wavelength photometric catalogs online and present approximate physical and evolutionary parameters for all of our YSO candidates. Our primary modes of source analysis are CMD examination and SED fitting. Optical $\it HST$ CMDs reveal $\sim$565 PMS candidates, essentially low mass Stage~{\sc iii} YSOs, through isochrone fitting. Through multi-wavelength SED fitting, we identify 41 YSO candidates, including 24 Stage~{\sc i}, 8 Stage~{\sc i}/{\sc ii}, 5 Stage~{\sc ii}, 2 Stage~{\sc ii}/{\sc iii}, and 2 unclassified candidates. High-resoultion {\it HST} imaging shows that $\ga 70\%$ of the YSO candidates include multiple sources or are protoclusters, and most of these optical sources are PMS candidates. Efforts to construct YSO protocluster models and incorporate them into the SED fitter are underway but are beyond the scope of this paper. For the $\sim$20 YSO candidates, we are able to estimate masses for one or more optical counter-parts via comparison with CMD evolutionary tracks, and we find consistency between these lower limit optical masses and the YSO SED fitter masses. We also construct a mass function from YSO SED fitter masses and derive a present-day star formation rate of 0.2-1~$\Msun$~yr$^{-1}$~kpc$^{-2}$. Finally, we present a quantitative analysis of the spatial distribution of the YSO population with respect to the central cluster and PMS population. We find that star formation has progressed from cluster center to the edge of the star forming dust cloud in NGC~602. The PMS stars are heavily concentrated near cluster center and that the YSO population distribution can be represented as concentric shells with Stage~{\sc ii} sources preferentially closer to cluster center and Stage~{\sc i} sources farther away. Previous observations \citep{nigra08} have shown that there is no significant expansion of the dust shell in which most of our YSO candidates are located and that the photo-ionization front is the prime mover of the star formation activity. We therefore correlate average distances of the Stage~{\sc i} and {\sc ii} YSOs from cluster center with the times at which their formation is triggered; we divide the distances by the sound speed. Relating the timescales, we find the lifetimes of each YSO Stage to be a few $10^5$~yr, comparable to timescale estimates in the literature which apply independent techniques to galactic star-formation regions.

1012.3406

1012

1012.1773_arXiv.txt

We compare the Gibbs and Maxwell constructions for the hadron-quark phase transition in neutron and protoneutron stars, including interacting hyperons in the confined phase. We find that the hyperon populations are suppressed and that neutrino trapping shifts the onset of the phase transition. The effects on the (proto)neutron star maximum mass are explored.

Some recent lattice QCD calculations have improved our knowledge on the nature of the nucleon-hyperon (NY) and hyperon-hyperon (YY) interactions \cite{nemura09}. Moreover, an intense experimental activity aimed at exploring hypernuclei has started at the J-PARC facility, which will hopefully clarify in the near future some unresolved questions on the hyperon interactions. This will help to improve the theoretical modeling of neutron star (NS) interiors, where hyperons are predicted to be present in large fractions, and many aspects of the equation of state (EOS) could be clarified. There are essentially two microscopic approaches to derive the properties and the EOS of dense matter starting from the bare baryon-baryon interactions, namely the Brueckner-Hartree-Fock (BHF) theory \cite{book} and the variational method \cite{apr}. Including the hyperon degrees of freedom, the EOS becomes very soft, and in the BHF approach the maximum mass for a NS lies below currently observed values of about 1.7 $\ms$ \cite{obs}. To overcome this problem, some possibilities have been suggested: i) The repulsive effects of three-body forces (TBF) among nucleons and hyperons, which may stiffen the EOS \cite{nishizaki02}; ii) Additional repulsion coming from hyperon-hyperon interactions (but no experimental data are available so far); iii) The hadron-quark phase transition \cite{burgio02,maieron04}, which could lead to a partial suppression of the hyperon population, thus stiffening the EOS. This conclusion needs to be further explored, since at present there are many uncertainties regarding the quark matter EOS. The scope of this work is to present results on the composition and structure of neutron and protoneutron stars (PNS), which are produced in the aftermath of successful supernova (SN) explosions of very massive stars ($M\gtrsim8\;\ms$). These objects can reach temperatures of the order of 30--40 MeV in their interiors, and are characterized by a temporary neutrino trapping with a conserved lepton fraction $Y_e \approx$ 0.4 \cite{burr,pons}, lasting some seconds. Both thermal effects and neutrino trapping may result in observable consequences of the neutrino signature from a supernova, and may also play an important role in determining whether a SN ultimately produces a cold NS or a black hole. In this paper, we focus on the hadron-quark phase transition, with the explicit inclusion of interacting hyperons in the confined phase. For hadronic matter we adopt the BHF microscopic EOS extended to finite temperature, and use the MIT bag model in the quark phase. This paper is organized as follows. In Section II we illustrate the Brueckner-Bethe-Goldstone (BBG) many-body theory including hyperons at finite temperature. Sections III and IV briefly review the MIT bag model and the treatment of the hadron-quark mixed phase, respectively. In Section V we discuss our results regarding the structure of NSs and PNSs, in particular their maximum mass. Final conclusions are drawn in Section VI.

We have studied the quark-hadron mixed phase in cold neutron stars and hot neutrino-trapped protoneutron stars containing also hyperons, and compared explicitly the Maxwell and the Gibbs phase transition constructions. We find that pure quark matter appears in the cores of compact stars in any situation and that the hyperon fractions are nearly completely suppressed by the appearance of quarks. Due to this reason the maximum NS mass is increased by the presence of quark matter. However, the simple MIT bag model used here is not capable to reach currently observed NS masses of about 1.7 $\ms$ \cite{obs}, and it will be an important task for the future to implement more sophisticated quark models. Apart from this challenge, in this paper we did not consider finite-size effects in the mixed phase. In particular, it will be interesting to check the influence of trapped neutrinos on the pasta structures in future works.

1012.1773

1012

1012.1897_arXiv.txt

% We present observations of a chromospheric jet and growing ``loop" system that show new evidence of a fan-spine topology resulting from magnetic flux emergence. This event, occurring in an equatorial coronal hole on 2007 February 9, was observed by the \hinode Solar Optical Telescope in unprecedented detail. The predecessor of the jet is a bundle of fine material threads % that extend above the chromosphere and appear to rotate about the bundle axis at $\sim$$50 \kmps$ (period $\lesssim$$200 \s$). These rotations or transverse oscillations propagate upward at velocities up to $786 \kmps$. % The bundle first slowly and then rapidly swings up, with the transition occurring at the onset of an A4.9 flare. A loop expands simultaneously in these two phases (velocity: $16$--$135 \kmps$). % Near the peak of the flare, the loop appears to rupture; simultaneous upward ejecta and mass downflows faster than free-fall appear in one of the loop legs. The material bundle then swings back in a whiplike manner and develops into a collimated jet, which is orientated along the inferred open field lines with transverse oscillations continuing at slower rates. Some material falls back along smooth streamlines, showing no more % oscillations. At low altitudes, the streamlines bifurcate at presumably a magnetic null point and bypass an inferred dome, depicting an inverted-Y geometry. These streamlines closely match in space the late Ca loop % and X-ray flare loop. These observations are consistent with the model that flux emergence in an open-field region leads to magnetic reconnection, forming a jet and fan-spine topology. We propose that the material bundle and collimated jet represent the {\it outer spine} in quasi-static and eruptive stages, respectively, % and the growing loop is a 2D projection of the 3D {\it fan surface}.

\label{sect_intro} Due to magnetic buoyancy \citep{ParkerE.mag-buoyancy.1955ApJ...121..491P}, magnetic flux ropes are expected to emerge from the convection zone into the corona through the photosphere and chromosphere. Such emerging flux regions \citep[EFRs;][]{Waldmeier.1st-flux-emerg.1937ZA.....14...91W, Ellison.1st-flux-emerg.1944MNRAS.104...22E} give birth to sunspots and active regions \citep{WeartZirin.AR-birth.1969PASP...81..270W}. When observed on the solar disk, an EFR is usually seen as a new bipole in magnetograms that grows in size and magnetic flux \citep{ZwaanC.flux-emerg.1978SoPh...60..213Z}. The opposite polarities of the bipole separate from each other at typical velocities $\lesssim$$1\kmps$ \citep{HarveyK.MartinSF.EFR-separa-vel.1973SoPh...32..389H, ChouWang.EFR-separ-vel.1987SoPh..110...81C}. Upflows of $\sim$$1\kmps$ at the photospheric level were observed \citep{BrantsJ.EFR-blueshift-1km/s.1985SoPh...98..197B} and confirmed in recent MHD simulations \citep{Archontis.emergeMHD.2004A&A...426.1047A, Martinez-Sykora_MHDemergenceII.2009ApJ...702..129M}. Once in the low-$\beta$ corona, driven by its magnetic pressure, the emerging flux expands at relatively larger velocities of 10--$20 \kmps$, as observed in rising \Ha arch filaments \citep{BruzekA.arch-filament.1967SoPh....2..451B, ChouZirin.arch-filament-rise.1988ApJ...333..420C} and extreme ultraviolet (EUV) loops \citep{Yashiro.Shibata.TRACE-EFR.2000ASPC..205..133Y} in EFRs. Dense photospheric or chromospheric material dredged up by the emerging flux was found to consequently drain down the legs of arch filaments at 30--$50 \kmps$ \citep{BruzekA.arch-filam-drain.1969SoPh....8...29B, Roberts.arch-filam-drain.1970PhDT........19R}. Such velocities of rise and drainage have been reproduced in MHD simulations \citep{Archontis.emergeMHD.2004A&A...426.1047A, FanY.rot-sunspot.2009ApJ...697.1529F}. When a flux rope emerges into an open-field region (e.g., coronal hole), magnetic reconnection between the emerging and ambient fields is expected to take place, producing a flare and material ejection \citep{HeyvaertsJ.flux-emerg-flare-model.1977ApJ...216..123H}. Such ejections were observed as {\it surges} in \Ha \citep{NewtonH.surge-discovery.1934MNRAS..94..472N, Roy.surge1973PhDT.........7R, Kurokawa.Kawai.surge.flux-emerg1993ASPC...46..507K} and as {\it jets} at other wavelengths, including white light \citep{WangYM.WL.EUV.jet1998ApJ...508..899W}, UV \citep{Brueckner.Bartoe.UVjet1983ApJ...272..329B}, EUV \citep{Alexander.Fletcher.jet.1999SoPh..190..167A}, and soft X-rays \citep{ShibataK.1st-SXT-jet.1992PASJ...44L.173S, StrongK.SXT-jet1992PASJ...44L.161S}. A classification of standard and blowout jets was proposed \citep{Moore.jet-dichotomy.2010ApJ...720..757M}. Torsional motions or helical features found in surges or jets \citep{XuAA.surge-rotate.1984AcASn..25..119X, Kurokawa.untwist-filamt1987SoPh..108..251K, ShimojoM.jet-stat.1996PASJ...48..123S, Patsourakos.EUVI-jet2008ApJ...680L..73P} were interpreted as relaxation of twists from the emerging flux \citep{Shibata.Uchida.helic-jet.1986SoPh..103..299S, Canfield.surge-jet1996ApJ...464.1016C, Jibben.Canfield.twist-surge-stat2004ApJ...610.1129J}. Numerical simulations have been extensively employed to explained various aspects of solar jets \citep{Shibata.Uchida.helic-jetMHD.1985PASJ...37...31S, YokoyamShibata.jetModel1995Natur.375...42Y, GalsgaardK.jet-emerg.2005ApJ...618L.153G, Nishizuka.giantCaHjet.2008ApJ...683L..83N, Ding_MHD-emerg-jet.2010A&A...510A.111D}. The simplest end state of flux emergence into to a (locally) unipolar region is a {\it fan-spine} configuration (\citealt{LauYT.Finn.3D-null.1990ApJ...350..672L, Torok.fan-spine.twist-emerg.2009ApJ...704..485T}; cf., multiple nulls connected by separators, \citealt{Maclean.null-emergence.2009SoPh..260..299M}). As shown in Figure~\ref{fan-spine.eps}, consider a sufficiently small bipole emerging into a region of a larger scale that has a net, say, negative, flux. The emerged flux introduces two new patches of opposite polarities, with the positive patch ending up as a minority-polarity isolated in a negative polarity all around. Regardless of reconnection development, % none of the field lines from this minority patch can leave this region, because the larger, surrounding region has a net opposite flux, and thus these field lines must fountain back to the nearby photosphere. Immediately outside of this closed field is the open field, and a {\it dome} or {\it separatrix fan surface} lies in between. A magnetic {\it null point} is located on the top of this dome. A special open field line, the {\it separatrix spine}, passes through this null point to continue on into the dome interior and be rooted at the base of the atmosphere. The two parts of this spine, called the outer and inner, are identified in Figure~\ref{fan-spine.eps}. A fan-spine topology bears significant implications for solar eruptions \citep{AntiochosS.breakout.1998ApJ...502L.181A} and its signatures were found in anemone-like active regions \citep{AsaiA.anemone-AR.2008ApJ...673.1188A}, Eiffel tower shaped X-ray jets \citep{ShimojoM.jet-stat.1996PASJ...48..123S}, saddle-like loop structures \citep{FilippovB.3D-null.1999SoPh..185..297F}, and circular flare ribbons \citep{MassonS.oval-ribbon-spine-fan.2009ApJ...700..559M}. Its presence in coronal jets or flares has been confirmed by magnetic field extrapolations \citep{Fletcher.3D-reconn.2001ApJ...554..451F, Moreno-Insertis.EISjet2008ApJ...673L.211M} and widely reproduced in 2D or 3D MHD simulations \citep[e.g.,][]{YokoyamShibata.jetModel1996PASJ...48..353Y, PariatE.twist-jet-homologous.2010ApJ...714.1762P}. \begin{figure}[thbp] % \epsscale{0.7} % \plotone{f1a.eps} \plotone{f1b.eps} \caption[]{A fan-spine topology resulting from emergence of a bipole into a unipolar region: (a) 2D vertical cut; (b) 3D bird's eye view. The hatched region in (a) represents postulated bright emission (see Section~\ref{sect_discuss}). } \label{fan-spine.eps} \end{figure} The launch of the \hinode mission \citep{Kosugi.Hinode2007SoPh..243....3K} has offered new opportunities to study the relationship and underlying physics of flux emergence, jets, and fan-spine topology in unprecedented detail \citep[e.g.;][]{LiHui.Hinode-flux-emerg.2007PASJ...59S.643L, Shibata.CaHjet.2007Sci...318.1591S, Okamoto.emerg-helical.2008ApJ...673L.215O, % Morita_jet-SOT-Hida-Ca-spec2010PASJ...62..901M}. % In an earlier Letter \citep[][hereafter Paper~I]{LiuW.CaJet1.2009ApJ...707L..37L}, we reported an intriguing chromospheric jet observed by \hinode on 2007 February 9 and focused on the fine structure kinematics of the jet itself. In this paper, we present a multiwavelength study of the entire event in greater detail. In Section~\ref{sect_obs}, we provide context observations, infer the unipolar magnetic environment, and investigate the associated flare. In Section~\ref{sect_sotanalys}, we pay special attention to the material bundle as the predecessor of the jet, the accompanying growing loop system, and the inverted-Y shaped geometry suggested by the streamlines of falling jet material. In Section~\ref{sect_discuss}, we propose that flux emergence in the unipolar region gives rise to the formation of a fan-spine topology, in which we identify the observed material bundle and jet as the outer spine and the growing loop as a 2D projection of the 3D fan surface. We conclude this paper in Section~\ref{sect_conclude} and provide supplementary information % in the Appendices. % \begin{figure}[thb] % \epsscale{0.9} % \plotone{f2.eps} \caption[]{YNAO \Ha and various \stereo Ahead EUVI images at different stages of the event. They are rotated by $90\degree$ counter-clockwise, such that the solar north is to the left. Note the dark absorption feature of the jet material in (d) and the bright falling material in (f). } \label{ha_stereo.eps} \end{figure} \begin{table}[bthp] % \scriptsize % \caption{Event milestones.} \tabcolsep 0.05in % \begin{tabular}{ll} \tableline \tableline 02:14--02:30 & earlier, brief surge-like activity \\ % 02:32 & bundle of material thread appears \\ % 02:44 & Ca loop and overarching SXR loop appear \\ 02:49:02 ($t_1$) & onset of flare, and of fast rise of material bundle \\ & and Ca loop \\ 02:50:32 ($t_2$) & end of Ca loop lateral expansion \\ 02:51:12 ($t_3$) & material bundle's lower end turns from vertical \\ & rise to horizontal drift; \\ & Ca loop ``ruptures" (apex undetectable) \\ 02:51:44 ($t_4$) & elbow appears in material bundle; \\ & northern Ca loop leg retreats downward \\ 02:52:24 & material bundle apex starts to sweep northward \\ 02:52:40 ($t_5$) & orientation angle of material bundle axis reaches \\ & maximum near flare peak; simultaneous upward \\ & ejecta and downflow in northern Ca loop leg \\ 02:55 & Ca loop leg and overarching SXR loop invisible \\ \tableline \end{tabular} \label{table_timeline} \end{table} %

% \label{sect_conclude} We have presented multiwavelength observations and detailed analysis of a chromospheric jet and its accompanying growing loop. This extends the study presented in \href{http://adsabs.harvard.edu/abs/2009ApJ...707L..37L}{Paper~I} which focused on the fine structure and kinematics of the jet itself. We summarize our new observations as follows. \begin{enumerate} \item % Potential field extrapolation indicates that this event occurs in an equatorial {\bf coronal hole} and as expected, the jet is closely aligned with the open field lines % (Figures~\ref{pfss_global.eps} and \ref{sot_Bfield.eps}). \item % The predecessor of the jet is a {\bf bundle of material threads} ($\lesssim$$1 \arcsec$ wide) % extending from the chromosphere into the corona. This bundle exhibits transverse {\bf sinusoidal oscillations} across its axis, whose velocities range from $47 \pm 9$ to $58 \pm 11 \kmps$, periods from $162 \pm 11$ to $197 \pm 35 \s$, and amplitudes from $1.5 \pm 0.3$ to $1.7 \pm 0.3 \Mm$. Such oscillations propagate upward at velocities as high as $v_{\rm ph} =786 \pm 30 \kmps$ (Figures~\ref{tslice_perp.eps} and \ref{perp_fit.eps}). We interpret these as evidence of propagating torsional MHD % waves. \item % The material bundle first slowly and then rapidly swings up, with the orientation angle of its central axis from the limb growing by $>$$50\degree$ in 10 minutes (Figure~\ref{jet_vs_time.eps}). The transition from the {\bf slow to fast swing} phase coincides with the onset of an A4.9 flare, which heats the plasma to $T= 12.2 \pm 0.6 \MK$. The bundle then swings back in a whiplike manner and develops into a collimated jet (Figure~\ref{mosaic.eps}), which continues to exhibit transverse oscillations (see \href{http://adsabs.harvard.edu/abs/2009ApJ...707L..37L}{Paper~I}), but at fractionally slower rates than the earlier bundle mentioned above. \item % A {\bf loop expands} simultaneously in these two phases. It attains a uniform vertical velocity of $16.2 \pm 0.4 \kmps$ during the {\it gradual phase} and reaches $135 \pm 4 \kmps$ at the end of the {\it acceleration phase} (Figure~\ref{jet_vs_time.eps}). The initial slow rise velocity is similar to those of emerging fluxes found in \Ha arch filaments and in MHD simulations. The lateral expansion is asymmetric and dominated by the northward displacement of the northern leg. \item % The loop appears to rupture or collapse near the peak of the flare and its apex becomes undetectable first. The northern leg of the loop {\bf retreats downward} and material drains down to the photosphere at accelerations ($a= (-4.8 \pm 0.7) g_\sun$) greater than free-fall. At the same time, some material is {\bf ejected upward} ($a=(11.4 \pm 3.1) g_\sun$) in the same leg (Figures~\ref{mosaic_NLeg.eps} and \ref{NLeg.eps}; Section~\ref{subsect_discuss}). \item % Some material falls back along streamlines in the original direction of ascent, showing no more transverse oscillations (\href{http://adsabs.harvard.edu/abs/2009ApJ...707L..37L}{Paper~I}). Most of the streamlines swerve around an inferred dome extending above the chromosphere and characterized with a null point at its top, depicting an {\bf inverted-Y} geometry. These streamlines closely match in space the late Ca loop prior to its rupture, the X-ray flare loop, and the EUV absorption features (Figure~\ref{streamline.eps}). % \end{enumerate} We interpret (Section~\ref{sect_discuss}) these observations in the framework of the emergence of a twisted flux rope into an open field environment leading to the formation of a jet through magnetic reconnection \citep[e.g.,][]{HeyvaertsJ.flux-emerg-flare-model.1977ApJ...216..123H, YokoyamShibata.jetModel1995Natur.375...42Y, Moreno-Insertis.EISjet2008ApJ...673L.211M}, and the relaxation of twists transferred into the jet leading to its spin \citep{Shibata.Uchida.helic-jetMHD.1985PASJ...37...31S, Shibata.Uchida.helic-jet.1986SoPh..103..299S, Canfield.surge-jet1996ApJ...464.1016C, Torok.fan-spine.twist-emerg.2009ApJ...704..485T}. We further identify signatures of the {\bf fan-spine} topology throughout the event, from the precursor to post-eruption evolution. The {\it outer spine} is recognized as the {\it material bundle} that eventually develops into the {\it collimated jet}, while the {\it fan surface} is imaged as the {\it growing Ca loop} in projection. After the eruption, the presence of this magnetic skeleton is clearly implied by the streamline geometry of the falling material. Our observations and model share commonalities with their counterparts in the literature, while the major differences given by our new findings are as follows. \begin{enumerate} \item % The {\bf simultaneous} growth of the emerging flux and development of the resulting jet, synchronized in two stages (Figure~\ref{jet_vs_time.eps}), have been clearly established here for the first time, to the best of our knowledge. Growing loops (other than post-flare loops) were recently noted in X-ray jet events \citep{ShimojoM.XRT-fine-jet.2007PASJ...59S.745S, ChiforC.Hinode-jet2.2008A&A...491..279C}, but their detailed temporal evolution and relationship with the jet were not clear. \item % In previous models \citep[e.g.][]{YokoyamShibata.jetModel1995Natur.375...42Y}, reconnection between the emerging flux and the overlying field would immediately lead to the launch of an eruptive jet. In our case, when the reconnection rate is {\it moderate} early in the event, the jet, manifesting itself as the material bundle, undergoes {\it quasi-static} evolution for more than 20 minutes. We call this an ``{\bf intermediate jet}" stage, which later develops into the classical {\it eruptive} jet as a result of {\it fast} reconnection driven by the accelerating expansion of the emerging flux. These are analogous to the slow and fast reconnections in the energy-storage and -release stages of coronal jet simulations, respectively \citep{PariatE.twist-jet-homologous.2010ApJ...714.1762P}. \item % The whiplike motion of a jet has been predicted as a consequence of the sling-shot effect of the newly reconnected field lines \citep[e.g.,][]{Shibata.Uchida.helic-jetMHD.1985PASJ...37...31S, Canfield.surge-jet1996ApJ...464.1016C}, and has been observed as a {\it unidirectional} swing {\it away} from the accompanying flare where reconnection occurs. In our case, the axis of the material bundle {\bf swings back and forth}, and the previously predicted whip motion only applies to the second swing here when the material bundle moves into its collimated jet position. We interpret this as instabilities possibly % related to the catastrophic unload of excessive twists (Section~\ref{subsect_model}). \end{enumerate} A statistical study of similar \hinode events is required before more general conclusions can be drawn. The validity of our phenomenological interpretation shall also be rigorously checked against theoretical models, numerical simulations \citep[e.g.,][]{PariatE.twist-jet-homologous.2010ApJ...714.1762P}, and laboratory experiments \citep[e.g.,][]{BellanP.plasma-lab-jet.2005Ap&SS.298..203B}. Finally, as pointed out by \citet{AntiochosS.breakout.1998ApJ...502L.181A}, the emergence of a bipolar flux system in a unipolar region on the photosphere naturally produces a local minority-polarity region with a spine-fan helmet magnetic structure above it. Our observations clearly show such an event and have identified its observable dynamical characteristics, thus providing motivation for future investigation of the rich 3D magnetic topologies to be found in such structures.

1012.1897

1012

1012.2085.txt

Collision of the magnetic flux tubes in the Quiet Sun was proposed as one of the possible sources for the heating of the solar atmosphere (Furusawa and Sakai, 2000). The solar photosphere was observed using the New Solar Telescope ad Big Bear Solar Observatory. In TiO spectral line at 705.68 nm we approached resolution of $0.1''$. The horizontal plasma wave was observed spreading from the larger bright point. Shorty after this wave an increase in the oscillatory power appeared at the same location as the observed bright point. This behavior matches some of the results from the simulation of the collision of the two flux tubes with a weak current.

The heating of the solar atmosphere is not well understood. Through the years of the research various solutions for the heating of the atmosphere were suggested and tested. One of the suggestions for the possible energy supply is collision of the magnetic flux tubes in the Quiet Sun (Furusawa \& Sakai, 2000). Furusawa and Sakai model present these collisions as the source of the fast magneto acoustic waves. Those waves should form shocks higher up in the atmosphere and deposit the energy trough the dissipation of the shock front. \par The solar photosphere has ubiquitous magnetic field. This fact is confirmed with the observations (Orzoco Su\'{a}rez {\it et al.}, 2007; Lites {\it et al.}, 2008) and simulations (Sch\"{u}ssler and V\"{o}gler, 2008; Steiner {\it et al.} 2008). Because the magnetic field is ubiquitous the importance of its study increases in the the quiet Sun, as well. The New Solar Telescope (NST) at Big Bear Solar Observatory (BBSO) revealed previously unresolved small structures in the quiet Sun photosphere. Moreover, previously unresolved bright points turned out to be a source for the part of the oscillations that were previously exclusively contributed to the intergranular lanes (Andic {\it et al.} 2010). We are now able to resolve those small structures and study their dynamic in detail. \par With NST at BBSO we were able to observe bright points in such detail that we're able to detect some of the signature proposed for collision of the magnetic flux tubes.

The ring-like speed distribution around the bright point was followed closely in time by rise in oscillatory power detected at the same location at which was the bright point in question. This sequence of the events might indicate that there was a collision of the magnetic flux tubes happening inside that bright point. Due to the limitation of our data set this conclusion is speculative. Our data set consist only from the intensity images, without any magnetic information nor information about line of sight velocities. Hence, this result needs to be tested with more complete data-sets.\par Although we still cannot resolve the substructure of the bright point itself, this signature points to another speculation that inside of the individual BP are individual flux tubes, indicating that the individual BP might be composed from the smaller, unresolved flux tubes.

1012.2085

1012

1012.2347_arXiv.txt

Stellar Kinematic Groups are kinematical coherent groups of stars which may share a common origin. These groups spread through the Galaxy over time due to tidal effects caused by galactic rotation and disk heating, however the chemical information survives. The aim of chemical tagging is to show that abundances of every element in the analysis must be homogeneous between members. We have studied the case of the Hyades Supercluster in order to compile a reliable list of members (FGK stars) based on chemical tagging information and spectroscopic age determinations of this supercluster. This information has been derived from high-resolution echelle spectra obtained during our surveys of late-type stars. For a small subsample of the Hyades Supercluster, stellar atmospheric parameters ($T_{\rm eff}$, $\log{g}$, $\xi$ and [Fe/H]) have been determined using an automatic code which takes into account the sensibility of iron $EWs$ measured in the spectra. We have derived absolute abundances consistent with galactic abundance trends reported in previous studies. The chemical tagging method has been applied with a carefully differential abundance analysis of each candidate member of the Hyades Supercluster, using a well-known member of the Hyades cluster as reference. A preliminary research has allowed us to find out which stars are members based on their differential abundances and spectroscopic ages.

The sample was selected using kinematical criteria in $U$, $V$ galactic velocities taking a dispersion of approximately 10 km/s around the core velocity of the group (Montes et al. 2001). \\ We have also taken additional candidates and spectroscopic information about some of these stars from L\'opez-Santiago et al. (2010), Mart\'inez-Arn\'aiz et al. (2010), and Maldonado et al. (2010). Some exoplanet-host star candidates are also taken from Montes et al. (2010).

We have computed the stellar parameters and their uncertainties for 42 single main sequence Hyades Supercluster candidate stars, after that we have obtained the chemical abundances of 12 elements, and the differential abundances. From the chemical tagging analysis we have found that 27 stars from the original sample are homogeneous in abundances for all the elements we have considered (a 64 \% of the sample), 3 stars fail to be homogeneous in one element. A more detailed analysis to check the consistency between the different age indicators and the chemical homogeneity is in progress.

1012.2347

1012

1012.0945_arXiv.txt

We have observed 37 bright, polarized radio sources with the Allen Telescope Array (ATA) to present a novel analysis of their Faraday rotation properties. Each source was observed during the commissioning phase with 2 to 4 100-MHz bands at frequencies ranging from 1 to 2 GHz. These observations demonstrate how the continuous frequency coverage of the ATA's log-periodic receiver can be applied to the study of Faraday rotation measures (RMs). We use RM synthesis to show that wide-bandwidth data can find multiple RM components toward a single source. Roughly a quarter of the sources studied have extra RM components with high confidence (brighter than $\approx$\ 40 mJy), when observing with a RM resolution of roughly 100 rad m$^{-2}$. These extra components contribute 10\%--70\% of the total polarized flux. This is the first time multiple RM components have been identified in a large sample of point sources. For our observing configuration, these extra RM components bias the measurement of the peak RM by 10--15 rad m$^{-2}$; more generally, the peak RM cannot be determined more precisely than the RM beam size. Comparing our 1--2 GHz RM spectra to VLBA polarimetric maps shows both techniques can identify complicated Faraday structures in the sources. However, the RM values and fractional polarization are generally smaller at lower frequencies than in the higher-frequency VLBA maps. With a few exceptions, the RMs from this work are consistent with that of earlier, narrow-bandwidth, all-sky surveys. This work also describes the polarimetry calibration procedure and that on-axis ATA observations of linear polarization can be calibrated to an accuracy of 0.2\% of Stokes I. Future research directions include studying the time-dependent RM structure in Active Galactic Nuclei (AGNs) and enabling accurate, wide-area RM surveys to test models of Galactic and extragalactic magnetic fields.

Radio waves encode information not only about their origin, but about their entire path of propagation. One way this information is encoded is through Faraday rotation, the frequency-dependent rotation of the radiation polarization angle caused by dispersion in a magnetized plasma \citep{f44,b66}. Observationally, Faraday rotation is parameterized by the rotation measure (RM): \begin{equation} \label{dtdl} \rm{RM} = \Delta \theta/\Delta(\lambda^2), \end{equation} \noindent where $\theta$ is typically measured in radians and $\lambda$ in meters. The plasma dispersion law predicts RM induced by propagation along a line as: \begin{equation} \label{phi} \rm{RM} = 0.81 \int_{0}^{d} \! n_e \, \mathbf{B} \, \cdot d\mathbf{l} \, \rm{rad} \, \rm{m}^{-2}, \end{equation} \noindent where $n_e$ is in cm$^{-3}$ and $\mathbf{B}$ is in $\mu$G, and $d$\ is the distance to the source \citep{b66}. Thus, measurements of Faraday rotation constrain physical conditions critical to understanding a wide variety of problems. Measurements of RM have expanded our knowledge of magnetic fields in our own Galaxy and in other galaxies. Several efforts have been made to compile RM along lines of sight over large areas of the sky \citep{s81,b03}. Finding patterns in these RM values have been used to constrain Galactic magnetic structure and turbulence \citep{h89,h06,h08}. Similar studies of other galaxies constrain models for the amplification of galactic magnetic fields \citep{g05}. Observations of individual AGNs and the massive black hole in the Galactic center have been used to constrain the geometry of their magnetic fields and depolarization \citep{b99,z05}. Recently, Taylor et al.\ (2009; hereafter ``T09'')\defcitealias{t09}{T09}, have expanded the number of sources with measured RM by nearly two orders of magnitude. This was done by reanalyzing the 1.4 GHz NRAO VLA Sky Survey \citep[NVSS; ][]{c98}, producing RMs for 37,543 radio sources located throughout the sky north of declination --40\sdeg. Each source was observed in two bands, which, assuming Equation \ref{dtdl}, can be used to estimate RM. The high density and large coverage of the \citetalias{t09} sample make it very effective at statistically measuring the structure and strength of the magnetic fields. However, while the \citetalias{t09} RMs are undoubtedly precise, it is not clear that they are accurate. \citet{b66} first noted that the Faraday rotation does not require the polarization angle to change as $\lambda^2$. When the emitting and Faraday-rotating media are mixed or multiple sources with different RM are spatially unresolved, one observes complicated changes in the polarization angle with wavelength \citep{g84}. These kinds of changes, combined with the possibility of $n\pi$ ambiguities in measuring polarization angle, make it difficult to measure RM robustly. \citet{b05} noted that discrete sampling of the Stokes vector in $\lambda^2$-space constrains different kinds of Faraday structures. They introduce the concept of ``rotation measure synthesis'', which takes advantage of the mathematical similarity between how aperture synthesis is done with multiple antennas and how RM is measured by sampling multiple wavelengths. The output from this technique is a ``RM spectrum'' showing the amount of polarized brightness as a function of its Faraday rotation \citep{h09,mao10}. In parallel with this algorithmic development has been the technical development of wide-bandwidth radio receivers and digital signal processing. The Allen Telescope Array (ATA), a radio interferometer in northern California, is commissioning these and many other new technologies \citep{w09}. One strength of the ATA design is its log-periodic receiver, which gives it continuous access to frequencies from 0.5 to 10 GHz \citep{w10}. The array and receiver design are optimized for large surveys \citep{c10}, which could be very powerful in the study of cosmic magnetism. The coincidence of these new algorithms and technologies for wide-bandwidth polarimetry inspired us to conduct a commissioning survey with the ATA. The first goal of this survey is to test the polarimetry capabilities of the ATA telescope design. Second, comparing the RM measured by the ATA and by \citetalias{t09} can test for biases associated with narrow bandwidth and coarse frequency resolution of the NVSS data. If the \citetalias{t09} RM measurements are trustworthy, it will give confidence in their application to measuring large-scale RM structure. Finally, the wide bandwidth and fine frequency resolution of the ATA demonstrates the power of RM synthesis to study complicated RM spectra. The paper begins with a description of the observations in \S \ref{obs}. The data reduction, including a detailed description of the ATA polarimetry calibration process, is described in \S \ref{red}. The analysis of the RM spectra is given in \S \ref{ana}, and implications of this work are given in \S \ref{dis}.

\label{con} We have conducted a survey of 37 bright, polarized radio sources from 1--2 GHz with the ATA. Applying polarimetric calibration allows us to measure the full Stokes information over this frequency range. RM synthesis uses these wide-bandwidth spectra to create high-resolution RM spectra, sensitive to RM up to 90000 rad m$^{-2}$ with a resolution as good as 50 rad m$^{-2}$. Observations of 3C 286 show that these data can measure RM with a precision up to 0.6 rad m$^{-2}$. In the 37 fields observed, 42 sources were brightly polarized enough to perform RM synthesis. Twelve of the 42 sources have multiple RM components with greater than 7$\sigma$\ significance ($\approx40$\ mJy), for a total of 61 components. Comparing the RM values measured by the ATA at high and low RM resolution show that multi-component RM spectra can bias the measurement of the peak RM of a source. We show that the peak RM of our sources cannot be known to better than 10--15 rad m$^{-2}$ if observed at RM resolution larger than 600 rad m$^{-2}$. More generally, the peak RM can be measured no more accurately than the width of the RM beam. However, for typical applications the mean RM measured by low-resolution RM spectra is not affected by this bias. The RM spectra measured by the ATA confirm the RM results presented by \citetalias{t09} for sources with polarized flux densities greater than 200 mJy. There is no evidence for $n\pi$\ ambiguity problems and the narrow-bandwidth VLA observations are at least 95\% reliable for this sample of sources. VLBA maps with milliarcsecond resolution spatially resolve some components that we identify in RM spectra. This correspondence suggests that RM synthesis can detect blobs of gas in massive AGN jets, which ordinarily must be spatially resolved to be studied. Observations of the time and frequency dependence of Faraday thickness and polarization fraction in AGN can be compared to simulations to constrain physical models. We have shown that RM synthesis is useful because it quantifies the full complexity of the Faraday rotation process. Future work will exploit the relative ease of the technique to resolve the time-dependent RM structure in AGN. Improving calibration models will also enable wide-field, wide-bandwidth RM surveys with the ATA and other radio interferometers. Large, accurate RM samples will probe magnetic fields on Galactic and extragalactic scales.

1012.0945

1012

1012.0518_arXiv.txt

The super-Earth GJ1214b transits a nearby M dwarf that exhibits $1\%$ intrinsic variability in the near-infrared. Here, we analyze new observations to refine the physical properties of both the star and planet. We present three years of out-of-transit photometric monitoring of the stellar host GJ1214 from the MEarth Observatory and find the rotation period to be long, mostly likely an integer multiple of 53 days, suggesting low levels of magnetic activity and an old age for the system. We show such variability will not pose significant problems to ongoing studies of the planet's atmosphere with transmission spectroscopy. We analyze 2 high-precision transit light curves from ESO's Very Large Telescope along with 7 others from the MEarth and FLWO 1.2 meter telescopes, finding physical parameters for the planet that are consistent with previous work. The VLT light curves show tentative evidence for spot occultations during transit. Using two years of MEarth light curves, we place limits on additional transiting planets around GJ1214 with periods out to the habitable zone of the system. We also improve upon the previous photographic $V$-band estimate for the star, finding $V=14.71\pm 0.03$.

\label{introduction} The transiting exoplanet GJ1214b offers an unparalleled opportunity to explore the physical properties of super-Earth planets. With a mass ($M_p$ = 6.6 \mearth) and radius ($R_p$ = 2.7 \rearth) between those of Earth and Neptune, and a likely equilibrium temperature ($T_{eq} = 500K$) cooler than for most transiting planets, GJ1214b represents an intriguing new kind of world with no Solar System analog \citep{charbonneau.2009.stnls}. Given intrinsic degeneracies in the mass-radius diagram in this regime \citep{seager.2007.mrse, adams.2008.optamrsewma,rogers.2010.fqdeic}, the bulk composition of the planet cannot be uniquely determined from current measurements of the mass and radius alone. For example, \citet{rogers.2010.tpol1} can explain the observed mass and radius to within $1\sigma$ with any of three generic physical models: (i) a mini-Neptune that accreted and maintained a low-mass H/He layer from the primordial nebula, (ii) a superfluid water-world with a sublimating ${\rm H_2O}$ envelope, or (iii) a rocky planet with an H-dominated atmosphere formed by recent outgassing. { Detailed calculations of GJ1214b's thermal evolution by \citet{2010arXiv1010.0277N} favor a metal-enriched H/He/\hho~envelope, finding that a water-only atmosphere would require an implausibly large water-to-rock ratio in the planet's interior.} Fortunately, because GJ1214b transits a very nearby (13 pc), bright ($K=8.8$), low-mass M dwarf (0.16 \msun), it is amenable to follow-up observations that could distinguish among these hypotheses. In particular, the large ($D=1.4\%$) transit depth enables transit studies of the planet's atmosphere. \citet{miller-ricci.2009.assdbhha} show that measuring the amplitude of the planet's transmission spectrum (i.e., the wavelength-dependence of the transit depth $\Delta D(\lambda)$ caused by absorption at the limb of the planet) constrains the mean molecular weight of its atmosphere and, in turn, the hydrogen content of its outer envelope. Cases (i) or (iii) of \citet{rogers.2010.tpol1} would produce $\Delta D(\lambda)\approx0.1\%$ variations in the transit depth across wavelengths accessible from the ground as well as {\em Hubble} and {\em Spitzer Space Telescopes}, while case (ii)'s dense atmosphere would result in variations below the sensitivity of current instruments \citep{miller-ricci.2010.nats1}. Providing a potential complication, however, the host star GJ1214 shows roughly sinusoidal photometric modulations that are presumably due to an asymmetric distribution of spots on a rotating star. Such spots can bias planetary parameters as measured from transit light curves whether or not they are occulted by the planet \citep[e.g.][]{pont.2007.hsttppt1mrs, desert.2011.tse1isom}, partially decoupling the observed transit depth $D$ from the actual planet-to-star radius ratio $R_p/R_{\star}$. Of particular importance for transmission spectroscopy studies, the change in transit depth induced by spots can vary with both time and wavelength, potentially mimicking the signal of a planetary atmosphere. Stellar spots have been observed in several transiting exoplanet systems around active stars \citep[see][]{strassmeier.2009.s}. {\em Hubble Space Telescope} ({\em HST}) photometry \citep{rabus.2009.csstpts} and later ground-based follow-up \citep{dittmann.2009.tdsdctepfgedtpsat} of TrES-1b has shown evidence for spot occultations in transit light curves. The high photometric precision and continuous coverage provided by the { CoRoT} satellite enabled detailed modeling of spotty transit and out-of-transit light curves for the hot Jupiters { CoRoT-2b} \citep{wolter.2009.tmscpsswpt, czesla.2009.saaseep, huber.2010.pemcedrsm} and { CoRoT-4b} \citep{aigrain.2009.npcdpp,lanza.2009.parpsc}. For the former, joint fits to the transit and out-of-transit flux showed that initial estimates of the planet's $R_p/R_{\star}$ were 3\% (9$\sigma$) too low \citep{czesla.2009.saaseep}. The interpretation of the transiting super-Earth { CoRoT-7b} is obfuscated by the fact that both the transit depth and the reflex motion are well below the amplitude of activity-induced modulations \citep{leger.2009.tefcsmvcfswmr,queloz.2009.cpsos}. Reanalyses of the { CoRoT-7} radial velocities find changing values for the mass of CoRoT-7b \citep{hatzes.2010.iirvvc,ferraz-mello.2010.pmdcsoasccs,lanza.2010.parrvvpsc} and call into question the significance of the mass measurements for both { CoRoT-7b} and the claimed outer planet { CoRoT-7c} \citep{pont.2010.rrepac}. Like GJ1214b, the well-studied hot Jupiter HD189733b \citep{bouchy.2005.emstjivjtbs1} is an ideal system for characterization studies, but requires corrections for stellar activity. The host HD189733 is an active K2 dwarf \citep{moutou.2007.sotpsd1} with 2\% peak-to-peak variability in the optical \citep{croll.2007.ls1sstmsp,miller-ricci.2008.msptes1ptmtaas}. \citet{henry.2008.rpps1} undertook a long-term photometric monitoring campaign from which they measured the 12 day stellar rotation period of HD189733. Extrapolation from their out-of-eclipse photometric spot characterization was useful for interpreting transmission spectroscopy results of individual transits from {\em Hubble} \citep[][]{pont.2007.hsttppt1mrs, swain.2008.pmaep} and measurements of the thermal phase curve from {\em Spitzer} \citep{knutson.2007.dcep1,knutson.2009.mcdcp1}. Understanding the time-variable surface of the star was even more crucial for broadband transmission spectroscopy studies that rely on comparing transit depths at different epochs \citep[e.g.,][]{desert.2009.scmate1,desert.2010.tsehisom,sing.2009.tse1iswfhwn}; interpretation of these data rely heavily on the photometric monitoring of \citet{henry.2008.rpps1}. To aid ongoing and future studies of GJ1214b, we present new data (\S\ref{Observations}) to characterize the star GJ1214's variability and estimate its rotation period (\S\ref{Rotation}). We compare the measured variability to a simultaneous analysis of 2 high-precision transit light curves from ESO's Very Large Telescope with 7 other new or previously published transits (\S\ref{Transits}). Additionally, we place upper limits on the radii of other possible transiting planets in the system (\S\ref{Search}) and present a refined estimate of the star's $V$ flux, which bears directly upon its metallicity as estimated using $M_K$ and $V-K$ relations. Finally, we discuss the implications of the measured variability for the properties of the star and for transmission spectroscopy studies of GJ1214b's atmosphere (\S\ref{Discussion}). We also note the following correction. In \citet{charbonneau.2009.stnls}, we quoted a systemic radial velocity for GJ1214 that had a typo in the sign; the actual velocity is $\gamma=+21.1\pm1.0$ km s$^{-1}$ (i.e. a redshift).

We have measured long-term photometric variability on GJ1214 to have a 1\% peak-to-peak amplitude { in the MEarth bandpass (715-1000 nm)} and a long rotation period, most likely an integer multiple of 53 days. Fitting very high precision light curves from the VLT, we find likely instances of GJ1214b crossing small spots during transit. Treating these occultation events as correlated noise, we find parameters for the planetary system that are consistent with previous work. We estimate the amplitude of time-variable changes in the apparent radius of the planet due to the observed stellar variability as $\Delta D_{\rm spots} (t) = 0.0001$ and place an upper limit of $\Delta D_{\rm spots}(\lambda) < 0.0003$ on possible spot-induced spectral features in the planet's transmission spectrum. Stellar spots do not limit current studies \citep[e.g.][]{bean.2010.temp, desert.2011.oemrasg}, but could be important for future studies of GJ1214b with {\em JWST}. Using two years of MEarth data, we have placed limits on the presence of other transiting planets around GJ1214. With 90\% confidence, we rule out the presence of Neptune-radius transiting planets in orbits shorter than 10 days but cannot place strong constraints on planets smaller than 2.0 \rearth~at such long periods. In a system where a 1.0\mearth~planet in a 2:1 mean motion resonance would create 100 second perturbations to GJ1214b's transit times, we find no evidence for transit timing variations larger than 15 seconds. Further searches of the GJ1214 system for potentially habitable planets smaller and cooler than GJ1214b continue to be warranted.

1012.0518

1012

1012.0338_arXiv.txt

Recent studies of precision electroweak observables have led to the conclusion that a fourth generation % is highly constrained. However, we point out that a long-lived fourth generation can reopen a large portion of the parameter space. In addition, it preserves baryon and lepton asymmetries against sphaleron erasure even if $B-L=0$. It opens up the possibility of exact $B-L$ symmetry and hence Dirac neutrinos. The fourth generation can be observed at the LHC with unique signatures of long-lived particles in the near future.

1012.0338

1012

1012.0424_arXiv.txt

Our understanding of the Universe today includes overwhelming observational evidence for the existence of an elusive form of matter that is generally referred to as dark. Although many theories have been developed to describe its nature, very little is actually known about its properties. Since its launch in 2008, the Large Area Telescope, onboard the Fermi Gamma-ray Space Telescope, has detected by far the greatest number ever of gamma rays, in the 20MeV 300GeV energy range and electrons + positrons in the 7 GeV- 1 TeV range. This impressive statistics allows one to perform a very sensitive indirect experimental search for dark matter. I will present the latest results on these searches. \vspace{1pc}

Recently the experimental information available on the Cosmic Ray Electron (CRE) spectrum has been dramatically expanded as the $Fermi$ LAT Collaboration~\cite{fermi} has reported a high precision measurement of the electron spectrum from 7 GeV to 1 TeV performed with its Large Area Telescope (LAT) \cite{Fermi_el}, \cite{Fermi_el2}. The spectrum shows no prominent spectral features and it is significantly harder than that inferred from several previous experiments. These data together with the PAMELA data on the rise above 10 GeV of the positron fraction \cite{Pam_pos} are quite difficult to explain with just secondary production~\cite{SM134},\cite{jcap}, \cite{IDM08}. \begin{figure}[p] \vskip -0.4 cm \includegraphics[width=8.2cm,height=7.8cm]{Fermi_el.eps} \vskip -0.9 cm \caption{\label{Fermi_el_fig} Cosmic ray electron + positron spectrum as measured by $Fermi$ Large Area Telescope for one year of observations (filled circles), along with other recent high energy results. The gray band represents systematic errors on the $Fermi$ LAT data \cite{Fermi_el}, \cite{Fermi_el2}. The solid line is the computed conventional GALPROP model but with an injection index $\Gamma $= 1.6/2.7 below/above 4 GeV (dotted line). An additional component with an injection index $ \Gamma = 1.5$ and exponential cut-off is shown by the dashed line. Blue line shows $e^-$ spectrum only. } \end{figure} \begin{figure}[!t] \centering \includegraphics[width=7.9cm,height=5.5cm ]{morselli_fig007.eps} \vskip -0.6 cm \caption{ The parameter space of dark matter particle mass versus pair-annihilation rate, for models where dark matter annihilates into monochromatic $e^\pm$ . Models inside the regions shaded in gray and cyan over-produce $e^\pm$ from dark matter annihilation with respect to the $Fermi$ LAT and H.E.S.S. measurements, at 2-$\sigma$ level. The red and blue contours outline the regions where the $\chi^2$ per degree of freedom for fits to the PAMELA and $Fermi$ LAT data is less than 1. } \label{fig:mod1} \vskip -0.6 cm \end{figure} The temptation to claim the discovery of dark matter from detection of electrons from annihilation of dark matter particles is strong but there are competing astrophysical sources, such as pulsars, that can give a strong flux of primary positrons and electrons (see \cite{puls0}, \cite{puls}, \cite{coutu}, \cite{int_pap} and references therein). At energies between 100 GeV and 1 TeV the electron flux reaching the Earth may be the sum of an almost homogeneous and isotropic component produced by Galactic supernova remnants and the local contribution of a few pulsars with the latter expected to contribute more and more significantly as the energy increases. \begin{figure}[ht] \hskip -0.6 cm \includegraphics[width=7.9cm,height=5.5cm]{ani.eps} \vskip -0.6 cm \caption{\label{anisotropy_f} Dipole anisotropy $\delta$ versus the minimum energy for GALPROP (solid line), Monogem source (dashed line), and Vela source (dotted line). The 95\% Upper limit's confidence level from the data is also shown with circles. The solar modulation was treated using the force-field approximation with modulation potential $\Phi$=550~MV. } \end{figure} \begin{figure}[ht] \includegraphics[width=7.9cm,height=7.2cm]{GC.4.eps} \vskip -0.5 cm \caption{ Counts spectra from the likelihood analysis of the $Fermi$ LAT data (number of counts vs reconstructed energy) in a 7$^{\circ} \times $7$^{\circ}$ region around the Galactic Center (number of counts vs reconstructed energy). } \label{GC_fig} \vskip -0.8 cm \end{figure} \begin{figure}[ht] \includegraphics[width=7.9cm,height=5.5cm]{res2.eps} \vskip -0.5 cm \caption{ Residuals ( (exp.data - model)/model) of the above likelihood analysis. The blue area shows the systematic errors on the effective area.} \label{GC_r_fig} \end{figure} Two pulsars, Monogem, at a distance of 290~pc and Geminga, at a distance of 160~pc, can give a significant contribution to the high energy electron and positron flux reaching the Earth and with a set of reasonable parameters of the model of electron production the $Fermi$ LAT data and the PAMELA positron fraction can be well fit fraction~\cite{Pam_pos} (see figure \ref{Fermi_el_fig}). However we have a lot of freedom in the choice of these parameters because we still do not know much about these processes, so further study on high energy emission from pulsars is needed in order to confirm or reject the pulsar hypothesis. Nevertheless a dark matter interpretation of the $Fermi$ LAT and of the PAMELA data is still an open possibility. Figure \ref{fig:mod1} shows the parameter space of dark matter particle mass versus pair-annihilation rate, for models where dark matter annihilates into monochromatic $e^\pm$ \cite{int_pap}. The preferred range for the dark matter mass lies between 400 GeV and 1-2 TeV, with larger masses increasingly constrained by the H.E.S.S. results \cite{HESS}. The required annihilation rates, when employing a particular dark matter density profile imply typical boost factors ranging between 20 and 100, when compared to the value $\langle\sigma v\rangle\sim3\times 10^{-26}\ {\rm cm}^3/{\rm sec}$ expected for a thermally produced dark matter particle relic. How can one distinguish between the contributions of pulsars and dark matter annihilations? Most likely, a confirmation of the dark matter signal will require a consistency between different experiments and new measurements of the reported excesses with large statistics. The observed excess in the positron fraction should be consistent with corresponding signals in absolute positron and electron fluxes in the PAMELA data and all lepton data collected by $Fermi$ LAT. $Fermi$ LAT has a large effective area and long projected lifetime, 5 years nominal with a 10 years goal, which makes it an excellent detector of cosmic-ray electrons up to $\sim$1 TeV. Future $Fermi$ LAT measurements of the total lepton flux with large statistics will enable distinguishing a gradual change in slope as opposed to a sharp cutoff with high confidence \cite{dark2}. The latter can be an indication in favor of the dark matter hypothesis. Another possibility is to look for anisotropies in the arrival directions of the electrons. The $Fermi$ LAT) detected more than 1.6 million cosmic-ray electrons/positrons with energies above 60~GeV during its first year of operation. The arrival directions of these events were searched for anisotropies of angular scale extending from $\sim$10$^\circ$ up to 90$^\circ$, and of minimum energy extending from 60~GeV up to 480~GeV. An upper limit for the dipole anisotropy has been set to 0.5 - 10\% depending on the energy \cite{anisotropy_p}. The levels of anisotropy expected for Vela-like and Monogem-like sources (i.e. sources with similar distances and ages) seem to be greater than the scale of anisotropies excluded by the results (see figure \ref{anisotropy_f} ). However, it is worth to point out that the model results are affected by large uncertainties related to the choice of the free parameters.

Fermi Gamma-ray Space Telescope has opened a new era in DM searches and a large variety of analyses have been developed for clusters of galaxies, DM satellites, DM subhalos, cosmological DM and spectral lines. No significant detections have been made, but constraints that start to probe the available phase space have been put on the annihilation cross-section and decay lifetimes. In addition, several ongoing analyses are now being finalized, including studies of the complicated galactic center region. The CRE spectrum measured by $Fermi$ LAT is significantly harder than what was expected on the basis of previous data. Adopting the presence of an extra $e^\pm$ primary component with $\sim$ 2.4 spectral index and $E_{cut} \sim 1$ TeV allows a consistent interpretation of the $Fermi$ LAT CRE data, HESS and PAMELA. Such an extra-component can be produced by nearby pulsars for a reasonable choice of relevant parameters or by annihilating dark matter for models with $M_{DM} \sim$ 1 TeV. Improved analysis and complementary observations (CRE anisotropy, spectrum and angular distribution of diffuse $\gamma$, DM sources search in $\gamma$) are required to possibly discriminate the right scenario. The dark matter origin of any exotic signal has to be confirmed by complementary findings in $\gamma$-rays by $Fermi$ LAT and atmospheric Cherenkov telescopes, and by LHC in the debris of high-energy proton destructions. On the other hand, if the signal is due to to be a conventional astrophysical source of cosmic rays, it will mean a direct detection of particles accelerated at an astrophysical source, again a major breakthrough. However, independent of the origin of these excesses, exotic or conventional, we can expect a very exciting several years ahead of us.

1012.0424

1012

1012.0321.txt

%% Text of abstract X-ray polarimetry promises to give qualitatively new information about high-energy sources. Examples of interesting source classes are binary black hole systems, rotation and accretion powered neutron stars, Microquasars, Active Galactic Nuclei and Gamma-Ray Bursts. Furthermore, X-ray polarimetry affords the possibility for testing fundamental physics, e.g.\ to observe signatures of light bending in the strong gravitational field of a black hole, to detect third order Quantum Electrodynamic effects in the magnetosphere of Magnetars, and to perform sensitive tests of Lorentz Invariance. In this paper we discuss scientific drivers of hard ($>$10 keV) X-ray polarimetry emphasizing how observations in the hard band can complement observations at lower energies (0.1 - 10 keV). Subsequently, we describe four different technical realizations of hard X-ray polarimeters suitable for small to medium sized space borne missions, and study their performance in the signal-dominated case based on Monte Carlo simulations. We end with confronting the instrument requirements for accomplishing the science goals with the capabilities of the four polarimeters.

\label{intro} Compared to observations in other parts of the electromagnetic spectrum, X-ray observations are of particular interest for the study of mass accreting black holes and neutron stars because these objects are X-ray bright and the X-rays originate very close to the compact objects. It is thus not surprising that X-rays are key to explore the properties of these objects. Whereas several X-ray imaging, spectroscopy, and timing missions have made spectacular discoveries over the last three decades \citep{Giac:79,True:83,Weiss:02,Stru:01,Jaho:96,Mi:07}, only one dedicated X-ray polarimetry mission has been launched so far. One of the reasons is that it is difficult to measure the polarization of an X-ray beam: whereas the arrival direction, arrival time, and energy of individual photons can be measured with extremely high accuracy, many hundreds of photons are needed to make even rough measurements of the three Stokes parameters $P$, $Q$, and $V$ which characterize the polarization properties of an X-ray beam. The only dedicated X-ray polarimetry mission to date OSO-8 \citep{Novi:75} detected a 2.6~keV and 5.2~keV polarization of the X-rays from the Crab Nebula of $\sim$20\% and a polarization angle aligned around 30 degrees oblique to the X-ray jet \citep{Weis:78}. For Cyg X-1 weak evidence for polarization on a level of a few percent was found \citep{Silv:79}; for other galactic compact objects upper limits on the polarization degree of a few 10 percent were measured \citep{Silv:79,Hugh:84}. Recently, two instruments on the {\it INTEGRAL} satellite were used to constrain the polarization of the hard X-ray emission from the Crab Nebula. Based on the analysis of data from the SPI instrument (SPectrometer on {\it INTEGRAL}), \citet{2008Sci...321.1183D} report tentative evidence for a 46\%$\pm$10\% polarization degree of the 100 keV-1~MeV emission. The analysis of 200 keV - 1 MeV data from the IBIS instrument (Imager on Board the {\it INTEGRAL} Satellite) indicates an even higher polarization fraction \citep{Foro:08}. The polarization direction seems to be aligned with the orientation of the X-ray jet at these energies \citep{2008Sci...321.1183D,Foro:08}. Models predict that galactic sources (e.g.\ binary black holes and neutron stars) and extragalactic sources (e.g.\ blazars, Gamma-Ray Bursts, GRBs) exhibit linear polarization degrees of a few percent and a few tens of percent, respectively, slightly below the sensitivity of the OSO 8 experiment. A mission with an order of magnitude improved sensitivity over OSO 8 should thus be able to detect the polarization of many objects, and to provide spectacular galactic and extragalactic results. Recent technological progress namely photo-electron tracking gas detectors \citep{Hill:07,Costa:08} have opened up the possibility to design small X-ray missions with more than two orders of magnitude better polarization sensitivities than OSO 8. The Gravity and Extreme Magnetism SMEX ({\it GEMS}) mission \citep{Swan:10} with excellent polarimetry sensitivity in the 2-10 keV energy band has recently been approved as a NASA mission. {\it GEMS} is projected to achieve a Minimum Detectable Polarization (MDP) degree of about 3\% for a mCrab source and an integration time of 1000 ksec. {\it GEMS} will have a single-photon energy resolution of between 15\% and 20\% Full Width Half Maximum (FWHM) and no imaging capabilities. At 0.5 keV {\it GEMS} will fly a student polarimeter with modest sensitivity. % We expect that {\it GEMS} will fulfill the high expectations and will motivate one or several follow-up missions. A follow-up mission may feature: \begin{itemize} \item improved sensitivity over the 2 keV - 10 keV energy range combined with an improved energy resolution, \item a broader energy bandpass with excellent sensitivity at lower ($<$2 keV) and/or higher ($>$ 10 keV) energies, \item the capability to do spectroscopic imaging polarimetry enabling the acquisition of 2-D maps of extended sources with spectroscopic and polarimetric information, \item a wide field of view (FoV) polarimeter with the possibility of measuring the polarization of transient sources as for example Gamma-Ray Bursts. \end{itemize} In this paper, we focus on the possibility to measure the polarization of hard X-rays ($>$10 keV) with narrow FoV instruments and wide FoV instruments. The soft gamma-ray telescope on board of the Japanese/US {\it ASTRO-H} mission (launch foreseen in 2013) will be able to do some hard X-ray Compton polarimetry using a combination of a collimator, Si pad detectors and CdTe pixel detectors \citep{Taka:08}. With an effective area of $>$30 cm$^2$ for Compton scattered events, the soft gamma-ray telescope on board of {\it ASTRO-H} should be able to verify some of the theoretical predictions discussed in this paper if on-ground and in-orbit calibration measurements can be used to reduce the systematic uncertainties below the level of the observed polarization effects. % In Section \ref{science} we summarize the science drivers for hard X-ray polarimetric observations. In Section \ref{designs} we discuss four different experimental approaches suitable for small to mid size space missions, and present a comparative study of the performance of the different polarimeters based on Monte Carlo simulations. In Section \ref{discussion} we summarize the results of the previous sections and critically discuss which science objectives may be addressed with the experimental approaches discussed before. The interested reader can consult \citep{Lei:97,Weis:06,Bell:10} for reviews of X-ray polarimetry and for information about different X-ray polarimeters. %

\label{discussion} Hard X-ray observations have to cope with the steep energy spectra of most astrophysical sources, and thus with lower photon fluxes than those in the soft X-ray band. However, hard X-ray polarimetry is an exciting upcoming field owing to several facts: \begin{itemize} % \item Several phenomena can {\it only} be observed at hard X-rays. Examples are the measurement of the high-energy end of the thermal emission from BBHs coming from the immediate surrounding of the black holes, observations of polarized cyclotron lines of magnetars, and the measurement of the polarization near the high-energy cutoffs of magnetars owing to the effect of photon splitting. % \item Almost all science topics that can be addressed with soft X-ray polarimetry benefit greatly from {\it simultaneous soft and hard X-ray spectropolarimetric observations}. The broadband energy dependence of the polarization degree and direction is crucial to verify that the models used to explain the soft X-ray polarization results are actually correct. In several cases combined soft and hard X-ray polarimetry observations are required to determine the model parameters that affect the interpretation of the results obtained in the two bands. A prominent example is the study of BBH systems: the combined soft and hard X-ray observations are needed to constrain the parameters describing the black hole, the accretion disk, and the corona. % \item In some sources the polarization degrees at higher energies are expected to be higher than at lower energies owing to the more compact emission regions of hard X-rays. For some sources the effect may make it easier to measure the hard X-ray polarization than the soft X-ray polarization. \item Hard X-rays allow us to study heavily obscured sources with column densities exceeding 10$^{24}$ cm$^{-2}$. \item For very hard sources (i.e. hard GRBs) hard X-ray observations achieve similar MDPs as soft X-ray observations as the photon number only depends logarithmically on the low energy threshold. \end{itemize} % \begin{table}[t] \begin{tabular}{p{5.2cm}|p{1.6cm}p{1.6cm}p{1.6cm}p{1.6cm}} \hline %\multicolumn{2}{|c|}{\small Trigger}\\ & Design 1 & Design 2 & Design 3 & Design 4\\ \hline \hline BBH thermal disk emission & \hspace*{0.6cm}\checkmark & \hspace*{0.6cm}- & \hspace*{0.6cm}- & \hspace*{0.6cm}- \\ BBH coronal emission & \hspace*{0.6cm}\checkmark & \hspace*{0.6cm}\checkmark & \hspace*{0.6cm}\checkmark & \hspace*{0.6cm}\checkmark\\ X-ray/$\gamma$-ray pulsars & \hspace*{0.6cm}- & \hspace*{0.6cm}\checkmark & \hspace*{0.6cm}\checkmark & \hspace*{0.6cm}\checkmark\\ NS cyclotron lines & \hspace*{0.6cm}\checkmark & \hspace*{0.6cm}- & \hspace*{0.6cm}- & \hspace*{0.6cm}- \\ NS vacuum birefringence & \hspace*{0.6cm}\checkmark & \hspace*{0.6cm}- & \hspace*{0.6cm}- & \hspace*{0.6cm}- \\ Magnetar X-ray tails & \hspace*{0.6cm}\checkmark & \hspace*{0.6cm}\checkmark & \hspace*{0.6cm}\checkmark & \hspace*{0.6cm} \checkmark\\ Magnetar photon splitting & \hspace*{0.6cm}- & \hspace*{0.6cm}\checkmark & \hspace*{0.6cm}\checkmark & \hspace*{0.6cm}\checkmark\\ AGN coronae & \hspace*{0.6cm}$\sim$ & \hspace*{0.6cm}- &\hspace*{0.6cm}-&\hspace*{0.6cm}-\\ Blazar jets & \hspace*{0.6cm}\checkmark & \hspace*{0.6cm}\checkmark & \hspace*{0.6cm}\checkmark & \hspace*{0.6cm}\checkmark\\ GRB jets & \hspace*{0.6cm}$\sim$$^a$ & \hspace*{0.6cm}\checkmark & \hspace*{0.6cm}\checkmark & \hspace*{0.6cm}\checkmark\\ Solar Flares & \hspace*{0.6cm}$\sim$$^b$ & \hspace*{0.6cm}\checkmark & \hspace*{0.6cm}\checkmark & \hspace*{0.6cm}\checkmark\\ Stellar Flares & \hspace*{0.6cm}- & \hspace*{0.6cm}- & \hspace*{0.6cm}- & \hspace*{0.6cm}-\\ LIV & \hspace*{0.6cm}\checkmark & \hspace*{0.6cm}\checkmark & \hspace*{0.6cm}\checkmark & \hspace*{0.6cm}\checkmark\\ \hline \end{tabular} \hspace*{1cm}$^a$ Requires alert by another instrument.\\ \hspace*{1cm}$^b$ A narrow field of view instrument would be unlikely to detect exceptionally strong solar flares without alerts from other instruments.\\ \caption{\label{sum} Science topics that can be addresses with the four different instrument designs.} \end{table} % % Table \ref{sum} shows the science topics that can be addresses with the different instrument designs. The table accounts for the different energy ranges of the four polarimeters and for their different field of views (Design 1: narrow field of view, Designs 2-4: possibly wide field of views). An estimate of the actual numbers of detected sources requires careful optimization of the shielding concept and is outside the scope of this paper. In the light of the results shown in the table three {\it GEMS} and {\it ASTRO-H} follow-up missions are attractive. The first mission is a narrow FoV broadband X-ray polarimetry mission with good sensitivity in the 0.1 keV-2 keV energy regime, an order of magnitude improved sensitivity and/or imaging spectropolarimetric capabilities in the 2 keV-10 keV energy band, and with spectropolarimetric coverage up to 80 keV. The second mission is a wide FoV observatory for spectropolarimetric studies of GRBs and flaring galactic sources, e.g.\ similar to EXIST, a large HX-POL, or POET. There is a niche for a large area detector assembly (similar to HX-POL, GRAPE) used with a pencil beam collimator to measure the $>$80 keV polarization properties of galactic sources, e.g.\ BBHs and the various flavors of neutron stars. The Washington University hard X-ray group is presently assembling a polarimeter called {\it X-Calibur} which adopts Design 1. We will report on detailed calibration measurements and on comparisons of simulated and experimentally measured data in a forthcoming paper. Pending approval by NASA, the polarimeter will be flown on a one-day balloon flight in the focal plane of an {\it InFOC$\mu$S} mirror assembly \citep{Ogas:05} with $\sim$40cm$^2$ detection area in spring 2012 and on subsequent longer balloon flights from Australia. In a collaboration with the Naval Research Laboratory (E.\ Wulff et al.), the group is also testing prototypes of the {\it HX-POL} Si-CZT polarimeter (Design 2). \\[2ex] % {\large \bf Acknowledgements:} The authors thank Martin Israel, Jonathan Katz (both Washington Univ.\ in St.\ Louis), Paolo Coppi (Yale), and J. Schnittman (Goddard Space Flight Center) for reading the manuscript carefully and for very valuable suggestions. Insightful comments by an anonymous referee helped to strengthen the paper substantially. HK acknowledges NASA for support from the APRA program under the grant NNX10AJ56G and support from the high-energy physics division of the DOE. The Washington University group is grateful for discretionary funding of the {\it X-Calibur} polarimeter by the McDonnell Center for the Space Sciences. % %% The Appendices part is started with the command \appendix; %% appendix sections are then done as normal sections %% \appendix %%

1012.0321

1012

1012.1600_arXiv.txt

The temperature fluctuations and polarization of the Cosmic Microwave Background (CMB) are now a well-known probe of the Universe at an infant age of 400,000 years. During the transit to us from the surface of last scattering, the CMB photons are expected to undergo modifications induced by the intervening large-scale structure. Among the expected secondary effects is the weak gravitational lensing of the CMB by the foreground dark matter distribution. We derive a quadratic estimator that uses the non-Gaussianities generated by the lensing effect at the four-point function level to extract the power spectrum of lensing potential fluctuations integrated out to $z \sim 1100$ with peak contributions from potential fluctuations at $z$ of 2 to 3. Using WMAP 7-year temperature maps, we report the first direct constraints of this lensing potential power spectrum and find that it has an amplitude of $A_L = 0.96 \pm 0.60$, $1.06 \pm 0.69$ and $0.97 \pm 0.47$ using the W, V and W+V bands, respectively.

Measurements of Cosmic Microwave Background (CMB) anisotropies have served as the strongest experimental probe of the early Universe to date (Spergel et al. 2003). Temperature fluctuations capture the physics of the primordial photon-baryon fluid undergoing oscillations in potential wells sourced by primordial density perturbations~(Hu et al. 1997; Hu \& Dodelson 2002). The anisotropies in the CMB also contain information related to the Universe at late times, as CMB photons in transit to us encounter large-scale structure. The main mechanisms that affect the frequency and direction of propagation of CMB photons are the gravitational interaction with time varying potential perturbations~(Sachs \& Wolfe 1967; Rees \& Sciama 1968), and Compton scattering by electrons due to the reionization of the universe~(Sunyaev \& Zel'dovich 1970). These mechanisms create specific signatures in the CMB that can be used to extract properties of the large-scale structure. One observational signature present in the CMB is the gravitational lensing modification by the projected dark matter distribution integrated along the line of sight (Hu 2000; Lewis \& Challinor 2006). Thi is analogous to weak lensing of galaxy shapes, now understood to be a strong probe of the dark matter distribution~(Wittman et al. 2000; Bacon et al. 2000, van Waerbeke et al. 2000; Bartelmann \& Schneider 2001). Unlike the case of lensing measurements with galaxy shapes, which restrict studies out to the era of $z \sim 2$, lensing of the CMB traces all the way back to the surface of last scattering at $z \sim 1100$. The dominant contributions to CMB lensing arise from $z \sim2$, but with a 30\% contribution at higher redshifts (Lewis \& Challinor 2006). From the ability to probe distance ratios and the integrated matter power spectrum at early times, CMB lensing is understood to be a powerful probe of certain cosmological parameters such as the neutrino mass and early dark energy~(Kaplinghat et al. 2003, Smith et al. 2008, Lesgourgues et al. 2006, Li \& Cooray 2006). Unlike lensing measurements with discrete galaxies, the primordial CMB sky is a continuous field, and a different technique is therefore needed. Due to lensing deflections, the temperature anisotropy $\Theta({\bf \hat n}) \equiv \Delta T/T({\bf \hat n})$ measured on the sky becomes $\tilde \Theta({\bf \hat n}+{\bf \alpha})$, where ${\bf \alpha} = \nabla \phi({\bf \hat n})$ is the deflection angle due to lensing given by the angular gradient of the lensing potential. The lensing potential is the line of sight projection between us and the last scattering surface of the gravitational potential $\Phi(r,\bn r)$ in the Universe: \begin{equation} \phi(\bn) = -2 \int_0^{r_0}dr {d_A(r_0-r) \over d_A(r) d_A(r_0)}\Phi(r,\bn r), \end{equation} where $d_A$ is the comoving angular diameter distance and $r$ is the comoving conformal distance from the observer. Using this expression, we can consider the dominant correction to the CMB temperature as a term that couples the deflection angle $\nabla \phi(\bn)$ to the gradient of the temperature: \begin{equation} \label{eq:pert} \Theta({\bf \hat n}) = \tilde \Theta({\bf \hat n + \alpha}) \sim \tilde \Theta({\bf \hat n}) + \nabla_i \phi({\bf \hat n}) \nabla^i \tilde \Theta({\bf \hat n}) , \end{equation} where $\tilde \Theta({\bf \hat n})$ is the unlensed temperature map. The CMB lensing effect is a modification to the angular gradient of temperature on the sky. If the CMB were completely isotropic, lensing modifications would not leave a change as lensing conserves surface brightness. Further, the lensing signatures are at second-order in temperature leading to a distinct non-Gaussianity pattern at the four-point function level, or trispectrum in Fourier space~(Zaldarriaga 2000; Hu 2001b). Quadratic statistics can be devised to probe the gradient structure of the CMB temperature and to extract the projected line of sight dark matter density field~(Hu 2001a; Cooray \& Kesden 2003; Okamoto \& Hu 2003). Attempts have been made for an indirect signature of the lensing effect in all-sky maps by Wilkinson Microwave Anisotropy Probe (WMAP) three-year data by estimating the deflection field and then cross-correlating that with a foreground density tracer field, such as NRAO VLA Sky Survey (NVSS) and Sloan Digital Sky Survey (SDSS). The overall signal-to-noise ratio for the measurement is 2.5$\sigma$~(Hirata et al. 2008) to 3.4$\sigma$~(Smith et al. 2007). With observations over seven years now completed, WMAP maps have improved in sensitivity to the extent that a direct measurement of the lensing signal can be pursued. We first construct an optimized statistic that probes the non-Gaussianity pattern induced by lensing via the trispectrum of the CMB temperature. We measure this by four maps weighted differently and taking a power spectrum of the temperature squared field. In addition to the lensing signal, this power spectrum contains a term associated with the Gaussian sky. We combine the constraints in data with a large suite of simulations, involving both the Gaussian temperature maps without lensing as well as maps with lensing included. This {\it Letter} is organized as follows. In the next Section, we outline the new estimator and how to weight CMB maps to extract the lensing signal. In Section~\ref{sec:test} we show that our estimator works correctly for simulated data. In Section~\ref{sec:Analysis} we describe our analysis used to constrain $C_l^{\phi \phi}$, the detection of the lensing amplitude and describe null tests used as a sanity checks. In Section~\ref{sec:final} we present our results and conclude with a discussion on implications for the future.

\label{sec:final} Fig.~\ref{fig:5fit} shows the scale dependent constraints on the lensing power spectrum obtained from the ${\cal K}_l^{(2,2)}$ estimator. We find the lensing signal to be compatible with the fiducial expectation. To our knowledge, this result provides the first direct constraints on $C_l^{\phi \phi}$ using CMB data encoding information for the matter distribution of the universe back to $z\sim1100$. Furthermore, this measurement does not appear to be biased by instrumental effects since the null test described above, is compatible with zero. The results for measuring the lensing amplitude $A_L$ using the CosmoMC analysis described in Section~\ref{sec:Analysis} are summarized in Table~\ref{tab:results}. We first reproduce values for the cosmological parameters consistent with Komatsu et al. (2010) for the case with no constraints on $C_l^{\phi\phi}$ and $A_L = 1$. Next we carry out the same run, without a constraint on $C_l^{\phi\phi}$, this time allowing $A_L$ to vary and find that WMAP 7-year data alone is consistent with an unlensed universe with $A_L = 0.87 \pm 1.05$. Finally, we find $A_L = 0.96 \pm 0.60$ and $A_L = 1.05 \pm 0.69$ when our constraints on the lensed power spectrum using the W and V frequency bands respectively are added to the WMAP 7-year data and is $A_L = 0.97 \pm 0.47$ when the W and V frequency bands are combined.. The amplitude of the $C_l^{\phi \phi}$ reported here is consistent with the naive expectation based on a Fisher matrix estimate that suggests a measurement no better than $2.5\sigma$. A variety of effects could be aiding the detection; our estimator is sensitive to the full trispectrum while a Fisher matrix estimate based on the reduced trispectrum with one pairing could underestimate the expected signal-to-noise ratio (see, Fig~3 of Hu 2001b). Effects such as the correlation between lensing and secondary anisotropies, such as the Sunyaev-Zel'dovich effect, could enhance the signal in the lensing trispectrum (Cooray \& Kesden 2003). While the full WMAP dataset will slightly improve the measurement we report here, a confirmation of our result showing a detection of the lensing power spectrum will come from Planck data, which is expected to make a larger than 60$\sigma$ detection of $A_L$ (Hu 2001b). As the measurement of $C_l^{\phi \phi}$ becomes more precise, it will provide tighter constraints on the cosmological parameters than can be obtained using temperature and polarization information alone. One obvious future target is a measurement of the sum of the neutrino masses, leading to a direct cosmological determination of the neutrino mass hierarchy~(Kaplinghat et al. 2003; Lesgourgues et al. 2006). A the lensing power spectrum further constrains the dark energy equation of state $w$ as well as early dark energy models in general (Kaplinghat et al. 2003; Joudaki in Prep.). The ultimate goal of CMB polarization measurements is the primordial gravitational wave signal in the so-called B-modes of polarization~(Baumann et al. 2009; Bock et al. 2009). The signal however is confused by the lensing effect that converts a small fraction of dominant scalar polarization in E-modes to B-modes~(Kesden et al. 2002; Knox \& Song 2002,Zaldarriaga, \& Seljak 1998). The fast estimator we have presented here for temperature maps can be generalized for CMB polarization maps (Munshi et al. in prep) and can be used to de-lens CMB B-modes to separate the signal from primordial gravitational waves and lensing of E-modes.

1012.1600

1012

1012.6020_arXiv.txt

In this paper we analyze high time resolution single pulse data of PSR B0809+74 at 820 MHz. We compare the subpulse phase behavior, undocumented at 820 MHz, with previously published results. The subpulse period changes over time and we measure a subpulse phase jump, when visible, that ranges from $95^\circ$ to $147^\circ$. We find a correlation between the subpulse modulation, subpulse phase, and orthogonal polarization modes. This variety of complicated behavior is not well understood and is not easily explained within the framework of existing models, most of which are founded on the drifting spark model of \citet{rud75}. We quantitatively fit our data with a non-radial oscillation model \citep{cle08} and show that the model can accurately reproduce the drifting subpulses, orthogonal polarization modes, subpulse phase jump, and can explain the correlation between all these features.

\label{intro} First discovered in 1968 by \citet{col68}, PSR B0809+74 is a bright, slow pulsar with drifting subpulses that has been continuously studied over the past 40 years. The literature contains a wide range of behavior including changes in subpulse period, subpulse phase behavior, average pulse shape, and orthogonal polarization modes as a function of radio frequency. The average pulse profile and polarization angle histogram for our observations at 820 MHz are shown in Figure \ref{fig:AverageProfile}. \citet{hob04} have monitored this pulsar for at least 6 years, measuring a spin period (\Pone) of 1.292 seconds and a dispersion measure of $6.116$ pc cm$^{-3}$. A list of basic parameters are in Table \ref{table:BasicParams}. \begin{figure} \begin{center} \includegraphics[scale=0.7]{f1.eps} \end{center} \caption{Top: The average total intensity (solid line) and linear polarization (dashed line) of PSR B0809+74 from epoch MJD 54922 at 820 MHz, consisting of 232 pulses. Bottom: A 2D histogram of the polarization angle for the same data. The center of the pulse profile is arbitrary; we chose maximum to be $60^\circ$ in pulse longitude to directly compare our data to that of \citet{edw03} (see Figure \ref{fig:edw03}). } \label{fig:AverageProfile} \end{figure} \include{tab1} At low frequencies, 81.5 to 151 MHz, the measured subpulse period is around 53 ms \citep{bar81,dav84}. The measured subpulse period is the spacing between two adjacent subpulses in the same pulse and is usually referred to as \Ptwo. As discussed in \citet{cle04}, the value of \Ptwo~ is not an accurate measurement of the underlying fundamental subpulse period, \Ptime. At higher frequencies, the \Ptwo~ appears to decrease to 39 ms, 31 ms, and 29 ms at 406 MHz, 1412 MHz, and 1720 MHz respectively \citep{dav84,bar81}. Our measurements at 820 MHz fall in the middle range of observational frequencies. \citet{bar81} find that while the subpulse period appears to change between 102.5 and 1720 MHz by a factor of 1.8, the time it takes for a subpulse to return to the same longitude, \Pthree, remains constant. The likely underlying cause for the change in the measured subpulse period is a subpulse phase discontinuity (or jump) that appears at high frequencies but not at low frequencies. The subpulse phase jump is not seen at 328 MHz \citep{edw03,edw04}, 408 MHz \citep{pro86}, or at 500 MHz \citep{wol81}. At 1380 MHz, \citet{edw03,edw04} report a phase jump of $\sim120^\circ$ as shown in the middle panel of Figure \ref{fig:edw03}. Furthermode, the subpulse phase jump occurs starts at approximately $56.5^\circ$ in pulse longitude (bottom left panel of Figure \ref{fig:edw03}), corresponding with a minimum the subpulse amplitude envelope (dark line in the top left panel of Figure \ref{fig:edw03}). As \citet{edw03} discuss, in any given pulse with two subpulses, the subpulses generally lie on opposite sides of the subpulse phase jump and appear closer together in pulse phase than they normally would in the absence of the phase jump, resulting in a smaller value of \Ptwo. \begin{figure} \begin{center} \includegraphics[scale=0.6]{f2.eps} \end{center} \caption{Top left: The average profile and subpulse longitude envelope for PSR B0809+74 at 328 MHz (dashed lines) and 1380 MHz (solid lines). Top right: The average profile and subpulse longitude envelope at 820 MHz taken on MJD 54922. Middle left: The subpulse phase, plotted multiple times, spaced $360^\circ$ apart. The white and dark circles are from data collected at 328 MHz and 1380 MHz, respectively. The dotted line shows the phase slope of $25^\circ$. Middle right: The subpulse phase for our data at 820 MHz; the dotted line shows the phase slope of $27.8^\circ$. Bottom left: The difference between the subpulse phase and that of the phase slope, indicating the magnitude of the phase jump. The phase jump at 1380 MHz (dark circles) is $120^\circ$ and the phase jump at 328 MHz (white circles) is plotted twice with an offset of $120^\circ$. Bottom right: The difference between the subpulse phase and that of the phase slope for our data at 820 MHz, results in a phase jump of $116^\circ$. Using an alternative method of fitting two lines on either side of the phase jump (middle, right panel) rather than subtracting the phase slope results in a phase jump of $145^\circ$ (see \S\ref{modes}). The phase jump at 820 MHz occurs at the same pulse longitude as the phase jump at 328 MHz and 1380 MHz. All the plots in the left panels are reproduced from \citet{edw03}.} \label{fig:edw03} \end{figure} Most drifting subpulse models are based on the drifting spark model \citep{rud75} in which a vacuum gap forms between the stellar surface and co-rotating magnetosphere due to the charge depletion from the emitted particles. To prevent the vacuum gap from growing indefinitely, sparks discharge across the vacuum gap. These sparks are fixed relative to each other and form a carousel that rotates around the magnetic pole at a rate incommensurate with the spin period of the star. The drifting subpulses are the manifestation of these spark discharges. Known as the drifting or (rotating) spark model, this model is the basis for many current models of drifitng subpulses \citep{kom70,bac76,gil00}. In \citet{cle04,cle08}, we proposed a non-radial oscillation model based on asteroseismological techniques \citep{dzi77} (see also \citet{rob82,cle00}) as an alternative to the drifting spark model. Pulsations in stars are not uncommon: white dwarf stars, ZZ Ceti stars, rapidly oscillating AP stars, and delta Scuti stars all show pulsation modes \citep{kle98,vank00,win81,kur82,bre69}. We were not the first proponents of a oscillation model for pulsars; \citet{gol68,vanh80,str92} all proposed oscillations as an explanation for drifting subpulses. However, these papers did not address the wide range of phenomenology seen in pulsars with drifting subpulses. In this paper, we analyze high quality single pulse measurements of PSR B0809+74 at 820 MHz. In \S\ref{obs} we discuss our observations and conduct a detailed analysis of the data in \S\ref{data}. We then explain our model in \S\ref{model} and examine the data in the context of our model in \S\ref{fit}. Finally, in \S\ref{conc}, we discuss how our 820 MHz observations compare the observations at other frequencies and explain the single pulse behavior within the context of a non-radial oscillation model.

This paper shows the second quantitative fit of our non-radial oscillation model to single pulse data of a pulsar \citep{ros08}. In this paper, we: \begin{itemize} \item show the subpulse period and subpulse phase jump, previously unpublished at 820 MHz. \item show the subpulse phase jump and subpulse period changes with epoch. \item are able to quantitatively determine the best value of \el~ and the subpulse period. Our method for determining the subpulse period is more accurate than using the FFT as it determines the subpulse period in real space rather than frequency space, which can be affected by the subpulse phase jump. Our fits using $I$, $Q$, and $U$ account for the subpulse phase jump since it is a natural part of the model. \item are able to quantitatively fit single pulse data to our model and determine a goodness of fit using \chisqr~ and ${\chi^2}_\nu$~ statistics. \item can create simulations based on our fitted parameters which accurately reproduce the subpulse period, subpulse phase jump, and orthogonal polarization modes. \end{itemize} The morphology of PSR B0809+74 is explained easily and naturally within a pulsational model. Our non-radial oscillation model is based on established asteroseismological principles that have explained white dwarf variations for the past 40 years \citep{dzi77}. This model is a viable alternative to the drifting spark model and can provide physical insight in the emission mechanism and physical structure of neutron stars.

1012.6020

1012

1012.2868_arXiv.txt

{\small We examine whether future, nearly all-sky galaxy redshift surveys, in combination with CMB priors, will be able to detect the signature of the cosmic neutrino background and determine the absolute neutrino mass scale. We also consider what constraints can be imposed on the effective number of neutrino species. In particular we consider two spectroscopic strategies in the near-IR, the so-called ``slitless'' and ``multi-slit'' approaches, whose examples are given by future space-based galaxy surveys, as EUCLID for the slitless case, or SPACE, JEDI, and possibly WFIRST in the future, for the multi-slit case. We find that, in combination with Planck, these galaxy probes will be able to detect at better than 3--sigma level and measure the mass of cosmic neutrinos: a) in a cosmology-independent way, if the sum of neutrino masses is above 0.1 eV; b) assuming spatial flatness and that dark energy is a cosmological constant, otherwise. We find that the sensitivity of such surveys is well suited to span the entire range of neutrino masses allowed by neutrino oscillation experiments, and to yield a clear detection of non-zero neutrino mass. The detection of the cosmic relic neutrino background with cosmological experiments will be a spectacular confirmation of our model for the early Universe and a window into one of the oldest relic components of our Universe.}

\label{Intro} Atmospheric and solar neutrino experiments have demonstrated that neutrinos have mass, implying a lower limit on the total neutrino mass given by $M_\nu\equiv \sum m_{\nu}\sim 0.05$ eV \cite{lesgourgues/pastor:2006}. This is a clear indication that the standard model for particle physics is incomplete and that there must be new physics beyond it. The neutrino mass splitting required to explain observations of neutrino oscillations indicates that two hierarchies in the mass spectrum are possible: two light states and a heavy one (normal hierarchy, NH, with $M_\nu>0.05$ eV), or two heavy and one light (inverted hierarchy IH, with $M_\nu>0.1$ eV). A third possibility is that the absolute mass scale is much larger than the mass splittings and therefore the mass hierarchy does not matter (degenerate neutrino mass spectrum). On-going and forthcoming neutrino experiments aim at determining the parameters of the neutrino mixing matrix and the nature of the neutrino mass (Dirac or Majorana). These experiments are sensitive to neutrino flavor and mixing angle, and to the absolute mass scale for large neutrino masses. As an example, beta-decay end-point spectra are sensitive to the neutrino mass, regardless of whether neutrinos are Dirac or Majaorana particles, and, the current limit on the effective electron neutrino mass is $< 2.2$ eV, coming from the Mainz and the Troitsk experiments, while KATRIN is expected to reach a sensitivity of $\sim 0.2$ eV \cite{Lobashev:2003kt,Kraus:2004zw,Thummler:2010tt}. Near future neutrino oscillation data may resolve the neutrino mass hierarchy if one of the still unknown parameters, which relates flavor with mass states, is not too small. However, if the mixing angle is too small, oscillation data may be unable to solve this issue. On the other hand cosmological probes are blind to flavor but sensitive to the absolute mass scale even for small neutrino masses (see Fig.1). In fact, a thermal neutrino relic component in the Universe impacts both the expansion history and the growth of structure. Neutrinos with mass $\la 1$ eV become non-relativistic after the epoch of recombination probed by the CMB, and this mechanism allows massive neutrinos to alter the matter-radiation equality for a fixed $\Omega_mh^2$. Neutrino's radiation-like behaviour at early times changes the expansion rate, shifting the peak positions in the CMB angular power spectrum, but this is somewhat degenerate with other cosmological parameters. WMAP7 alone constrains $M_\nu<1.3$ eV \cite{Komatsuetal2010} and, thanks to improved sensitivity to polarisation and to the angular power spectrum damping tail, forecasts for the Planck satellite alone give $M_\nu\sim 0.2-0.4$ eV, depending on the assumed cosmological model and fiducial neutrino mass (e.g., \cite{Perotto, Kitching_nu} and references therein). Massive neutrinos modify structure formation on scales $k > k_{\rm nr}=0.018(m_\nu/1{\rm eV})^{1/2}\Omega_m^{1/2}h$/Mpc, where $k_{\rm nr}$ is the wave-number corresponding to the Hubble horizon size at the epoch $z_{\rm nr}$, when a given neutrino species becomes non-relativistic. In particular, neutrinos free-stream and damp the galaxy power spectrum on scales $k$ larger than the so called free-streaming scale $k_{\rm fs}(z)=0.82 H(z)/(1+z)^2 (m_\nu/1{\rm eV}) h{\rm Mpc}^{-1}$ \cite{lesgourgues/pastor:2006}, thereby modifying the shape of the matter power spectrum in a redshift-dependent manner (see Fig.~\ref{transfer} and e.g. \cite{HuEisensteinTegmark,0709.0253,1004.4105,1003.2422}). Therefore, much more stringent constraints can be obtained by combining CMB data with large-scale structure (LSS) observations. Ref.~\cite{Reidnu,0505390} showed that present data-sets yield a robust upper limit of $M_\nu<0.3$ eV, almost ruling out the degenerate mass spectrum; this result was later confirmed by \cite{Lahavmassnu,Concha}. The forecasted sensitivity of future large-scale structure experiments, when combined with Planck CMB priors, indicate that cosmology should soon be able to detect signatures of the cosmic neutrino background and determine the sum of neutrino masses (e.g. \cite{HannestadWong,Hannestadreview,Kitching_nu,lsstbook,LahavDES} and references therein). Since cosmology is only weakly sensitive to the hierarchy \cite{1003.5918}, a total neutrino mass determination from cosmology will be able to determine the hierarchy only if the underlying model is normal hierarchy and $M_\nu<0.1$ eV (see e.g. Fig.~\ref{hierarchy}). \begin{figure*} \includegraphics[width=0.9\textwidth]{Hierarchy_definition.eps} \caption{Constraints from neutrino oscillations (shaded regions) and from cosmology. In this parametrisation the sign of the mass splitting specifies the hierarchy. The red triangles show the fiducial models explored in this work and the light blue vertical bands our forecasted errors (see \S 5). For fiducial $M_{\nu}$ values below $0.1$ eV a LCDM model must be assumed to obtain a detection with $> 2$--$\sigma$ statistical significance. For higher fiducial $M_{\nu}$, we can marginalise over dark energy parameters and still obtain tight errors on $M_{\nu}$.} \label{hierarchy} \end{figure*} A detection of the cosmic relic neutrino background (RNG) with cosmological experiments\footnote{Recall that neutrino experiments are not sensitive to relic neutrinos, as current generation of experiments do not have sufficient energy resolution to cleanly pin down the signature of the RNG. Anyway, the beta-decay end-point spectrum is in principle also sensitive to the RNG, and this can be foreseen as a plausible perspective for future experiments only if neutrinos have masses of order eV, thus in the so called degenerate scheme for neutrino masses, which is still allowed by all present data, though slightly disfavored by cosmological observations \cite{Cocco:2007za}.} would be a spectacular confirmation of our model for the early Universe and a window into one of the oldest relic components of our Universe besides the one represented by the stochastic gravitational wave background. This consideration prompts us to examine whether future galaxy redshift surveys probing LSS will be able to detect the signature of the neutrino background and to determine the neutrino absolute mass scale. Beyond neutrino mass, cosmology is also sensitive to the number of neutrino species. In the standard model for particle physics there are three neutrinos; they decouple early in the cosmic history and then contribute to the relativistic energy density (i.e. as if they were radiation) with an effective number of neutrino species $N_{\rm eff}=3.046$ (e.g. \cite{lesgourgues/pastor:2006}) until they become non-relativistic. Cosmology is sensitive to the physical energy density of relativistic particles, which include photons and neutrinos: $\Omega_r=\Omega_{\gamma}+N_{\rm eff}\Omega_{\nu}$, where $\Omega_{\gamma}$ and $\Omega_{\nu}$ are the energy density in photons and in one active neutrino species, respectively. CMB observations have constrained exquisitely well $\Omega_{\gamma}$, thus constraints in $\Omega_{\rm r}$ can be used to study neutrino properties. Deviations from $N_{\rm eff}=3.046$ would indicate non-standard neutrino properties or additional effective relativistic species. While the motivation for considering deviations from the standard model in the form of extra neutrino species has now disappeared \cite{mena,miniboone,mangano}, departures from the standard $N_{\rm eff}$ value could arise from decay of dark-matter particles \cite{bonometto98,lopez98,hannestad98,kaplinghat01}, early quintessence \cite{bean}, or more exotic models \cite{unparticles}. Relativistic particles affect the CMB and the matter power spectrum in two ways: {\it a)} through their anisotropic stress \cite{trotta/melchiorri:2005,Komatsuetal2010}, and {\it b)} through their relativistic energy density which alters the epoch of matter radiation equality. The ratio of CMB peak heights constrains matter-radiation equality yielding a degeneracy between $N_{\rm eff}$ and $\Omega_m h^2$. This degeneracy can be lifted by adding either cosmic expansion history data \cite{deberanrdisneff,Figueroa,hzstern} or adding the large-scale shape of the matter power spectrum: the power spectrum turnover scale is also related to matter-radiation equality given by the parameter $\Gamma\sim \Omega_m h$ (note the different scaling with $h$ compared to the CMB constraint). LSS surveys can yield a measurement, at the same time, of both the cosmic expansion history (via the Baryon Acoustic Oscillations (BAO) signal), and the large scale turnover of the power spectrum. Present constraints are already competitive with nucleosynthesis constraints, and future data will offer the possibility to test consistency of the standard paradigm for the early Universe. In fact, nucleosynthesis constraints rely on physics describing the Universe when its energy scale was $T\sim$ MeV, while cosmological constraints rely on physics at $T\sim eV$. In this paper we forecast errors on the total neutrino mass $M_\nu$ and the effective number of relativistic species $N_{\rm eff}$ by combining Planck priors with data from future space-based galaxy redshift surveys in the near-IR. In particular, we consider two main survey strategies: \begin{itemize} \item The first approach is to use ``multi-slit'' spectroscopy aimed at observing a pure magnitude-limited sample of galaxies selected in the near-IR (e.g. in the H-band at 1.6 $\mu$m) with a limiting magnitude appropriate to cover the desired redshift range. Examples of this approach are given by instruments where the efficient multi-slit capability is provided by micro-shutter arrays (MSA) (e.g. JEDI\footnote{http://jedi.nhn.ou.edu/} \cite{Wang04,Crotts05,Cheng06}), or by digital micromirror devices (DMD) (e.g. SPACE \cite{Cimatti09} and possibly WFIRST\footnote{http://wfirst.gsfc.nasa.gov/} in the future). With the multi-slit approach, all galaxy types (from passive ellipticals to starbursts) are observed, typically at $0<z<2-3$, if the observations are done in the near-IR, and provided that the targets are randomly selected from the magnitude-limited galaxy sample. \item The second approach is based on slitless spectroscopy (e.g. Euclid\footnote{http://sci.esa.int/euclid} and JDEM\footnote{http://jdem.gsfc.nasa.gov/} \cite{Glazebrook05,Laureijs09,JDEM}) which, due to stronger sky background, is sensitive mostly to galaxies with emission lines (i.e. star-forming and AGN systems), and uses mainly H$\alpha$ as a redshift tracer if the observations are done in the near-IR to cover the redshift range $0.5<z<2$. \end{itemize} Forthcoming surveys will also have a weak gravitational lensing component, which will also be used to constrain neutrino properties (see e.g. \cite{Kitching_nu}). Here we concentrate on galaxy clustering as an independent and complementary probe. The rest of the paper is organised as follows. In \S~\ref{Fisher matrix approach} we review our method and the employed modelling. In \S~\ref{spec_methods} we report the characteristics of the galaxy surveys considered in this work, and in \S~\ref{Fiducial cosmologies} we describe the adopted fiducial models and the explored space of cosmological parameters. In \S~\ref{Results} we present our results on the forecasted errors on the neutrino mass and number of neutrino species, and final in \S~\ref{Conclusions} we draw our conclusions.

\label{Conclusions} In this work we have forecasted errors on the total neutrino mass $M_\nu$ and the effective number of relativistic species $N_{\rm eff}$, by combining Planck priors with future data from space-based spectroscopic galaxy redshift surveys in the near-IR. We have considered two survey strategies based on slitless and multi-slit spectroscopies. The assumed set of cosmological parameters is very general and takes into account a time-varying dark-energy equation of state, as well as a non-vanishing spatial curvature of the Universe. We exploited information from the galaxy power spectrum shape and BAO positions, marginalising over galaxy bias; thus our findings do not depend on bias measurement accuracy (as long as, on the large scales considered, bias is scale independent or its scale dependence is known), or modelling of the redshift dependence of bias \cite{LahavDES}. The 1--$\sigma$ errors are shown in Tables~\ref{slitless_errors}-\ref{dmd_errors}, and the correlation coefficients in Tables~\ref{slitless_corr}-\ref{dmd_corr}. In Figs.~\ref{fig_Mnu03_sigmas}-\ref{fig_Mnu005_comp} we show the joint 2-parameter confidence levels. Regarding $M_\nu$--errors, we find that the multi-slit spectroscopy is able to reduce the neutrino mass errors of about 20\%-30\% compared to the slitless spectroscopy, depending on the fiducial total neutrino mass, if LSS data alone are used. When Planck priors are added, the 1--$\sigma$ errors on $M_\nu$ are in the range $0.03-0.05$ eV, depending on the fiducial neutrino mass, with an average difference of 15\% between the two spectroscopic strategies, favouring the multi-slit spectroscopy. Moreover, depending on the fiducial $M_\nu$--value, the total CMB+LSS dark-energy FoM, with growth--information marginalised over, decreases only by $\sim 15\%-25\%$ with respect to the value obtained if neutrinos are assumed to be massless (or their mass is assumed to be perfectly known), meaning that the ``$P(k)$--method marginalised over growth--information'' is quite robust to assumptions about model cosmology when constraining the dark-energy equation of state. The situation is different when we include growth-information, since in this case the value of the dark-energy FoM decreases by a factor $\sim 2-3$ with respect to cosmologies that assume massless neutrinos. Considering the fiducial cosmology with $M_\nu|_{\rm fid}=0.05$ eV, in \S \ref{growth&FoG} we checked the stability of $M_\nu$--errors to the inclusion of growth--information and peculiar velocity uncertainties. We compared the following approaches: the ``full $P(k)$--method, marginalised over growth--information'', the ``full $P(k)$--method, with growth--information included'', and ``full $P(k)$--method, with FoG and growth--information included''. We found that $M_\nu$--errors are quite stable at $\sigma(M_\nu)=0.05$ eV, against the adopted method. This result is as expected, if we consider that, unlike dark energy parameters, $M_\nu$ affects the shape of the power spectrum via a redshift-dependent transfer function $T(k,z)$, which is sampled on a very large range of scales including the $P(k)$ turnover scale, therefore this effect dominates over the information extracted from measurements of $f_g\sigma_8$. Regarding $N_{\rm eff}$--errors, again we find that, compared to the slitless spectroscopy, the multi-slit spectroscopy is able to reduce the $N_{\rm eff}$--errors by $\sim$30\% when LSS alone are used. When Planck priors are added, we find $\sigma(N_{\rm eff})\sim 0.08$, with only a $6\%$ difference between the two spectroscopy strategies, again in favour of the multi-slit one. The total CMB+LSS dark-energy FoM decreases only by $\sim 5\%$ with respect to the value obtained holding $N_{\rm eff}$ fixed at its fiducial value, meaning that also in this case the ``$P(k)$--method marginalised over growth--information'' is not too sensitive to assumptions about model cosmology when constraining the dark-energy equation of state. Finally, in Table~\ref{summary} we summarise the dependence of the $M_\nu$-- and $N_{\rm eff}$--errors on the model cosmology, for the two spectroscopic strategies combined with Planck. We conclude that, if $M_\nu$ is $>0.1$ eV, these surveys will be able to determine the neutrino mass scale independently of the model cosmology assumed. If $M_\nu$ is $<0.1$ eV, the sum of neutrino masses, and in particular the minimum neutrino mass required by neutrino oscillations, can be measured in the context of a $\Lambda$CDM model. This means that future spectroscopic galaxy surveys, such as Euclid or SPACE, JEDI, and possibly WFIRST in the future, will be able to cover the entire parameter space for neutrino mass allowed by oscillations experiments Moreover, as summarised in Fig.~\ref{hierarchy}, they will be competitive with future 3D cosmic shear photometric surveys, which, in combination with Planck priors, will give similar constraints on $M_\nu$ and $N_{\rm eff}$ \cite{Kitching_nu}. Since, these two kinds of LSS probe are affected by different systematics, their constraints on neutrino masses and relativistic degrees of freedom will provide a consistency check of the two independent measurement methods. We conclude that future nearly all-sky spectroscopic galaxy surveys will detect the cosmic neutrino background at high statistical significance, and provide a measurement of the neutrino mass scale. This will provide an important confirmation of our model for the early Universe, and crucial insights into neutrino properties, highly complementary to future particle physics experiments.

1012.2868

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1012.3480_arXiv.txt

The detailed workings of the central engines of powerful quasars remain a mystery. This is primarily due to the fact that, at their cosmological distances, the inner regions of these quasars are spatially unresolvable. Reverberation mapping is now beginning to unlock the physics of the Broad Emission Line Region (BELR) in nearby, low-luminosity quasars, however it is still unknown whether this gas is dominated by virial motion, by outflows, or infall. The challenge is greater for more distant, powerful sources due to the very long response time of the BELR to changes in the continuum. We present a new technique for probing the kinematic properties of the BELR and accretion disk of high-z quasars using differential microlensing, and show how substantial information can be gained through a single observation of a strongly-lensed quasar using integral field spectroscopy. We apply this technique to GMOS IFU observations of the multiply-imaged quasar Q2237+0305, and find that the observed microlensing signature in the CIII] broad emission line favours gravitationally-dominated dynamics over an accelerating outflow.

Differential microlensing offers the promise of indirectly measuring the spatial size and kinematics of different emitting regions in quasars. In this paper, a new observation of differential microlensing between the continuum and different velocity components of the CIII] broad line emission line in Q2237+0305 is presented. These data are then used to constrain the kinematics of the broad line gas. The core regions of quasars are unresolved at any wavelength, and so physical understanding must be based on theoretical models constrained by their observed spectral properties. The gas responsible for the broad emission lines presents a particular challenge as there are relatively few measurements capable of even beginning to probe its kinematical structure. As such, there is still no agreement on even the simplest models of this structure. Beyond the obvious width of the emission lines, which indicate large Doppler velocities of tens of thousands of km/sec, most of our current understanding has come from reverberation mapping and from polarisation observations of emission lines. The most recent reverberation signatures have been interpreted as arising from outflowing, inflowing and virialised motions \citep{Denney09, Bentz09}, while the polarisation studies have indicated an outflowing helical motion in the emission line gas \citep{Young07}. Theoretical models have favoured outflowing winds (eg \citealt{MurrayChiang97}), but alternative points of view are still under active consideration (eg. \citealt{Gaskell09}). Differential microlensing provides an independent probe of the inner structure of quasars. In quasars that are subject to strong gravitational lensing, the BELR and potentially the accretion disk are resolved on the spatial scale of the fine magnification structure of the lens. Given a model for the lensing galaxy, observation of differential magnification between these components allows constraints to be placed on their sizes. The use of this technique was first discussed by \citet{nem88}, who specifically looked at the effect of a single low mass star on a range of kinematical models for the broad emission line region (BELR). A key result of this paper was that different parts of the emission lines were differentially magnified, depending on the kinematical model: in general the smaller spatial regions showed the greatest variations, as would be expected. \citet{SW90} further considered the problem, noting that, in the case of a macro-imaged quasar, the differences between the lines in different images can be used to test whether particular emitting regions are being significantly microlensed. In particular, these authors noted that Keplerian motions in the BELR would be much easier to detect than infall (or presumably outflow). A key improvement in the modelling compared to \citet{nem88} was the use of microlensing magnification patterns to model the possible statistical variations for each macro-image. \citet{SW90} also noted that differential microlensing would also effect the redshift measured from a particular line. More recently \citet{Abajas07} and \citet{LewisAbata04} both considered different signatures which might be induced in the observed structure of the broad emission lines. There are at least four differential microlensing experiments which could allow the measurement of the physical parameters of the quasar emission regions. Target-of-opportunity observations of a quasar crossing a caustic provide the cleanest imaging experiment \citep{wwtm}. The physical interpretation of a caustic crossing event is straightforward. However in order to trigger the target-of-opportunity, regular monitoring is required; and for a reasonable annual probability of observing a caustic crossing, more than ten objects would need to be monitored. Recently, monitoring data has been used to fit the size of the region emitting the quasar continuum using Bayesian Monte Carlo methods \citep{koch04, Eigenbrod08b}. These analyses rely on detailed modeling where a standard accretion disk is convolved with microlensing networks. A third and more specialized possibility arises when two macro-images straddle a caustic. In a surprising number of cases, the fluxes of the two images differ, while theory predicts them to be the same. Several variables in the modeling can affect the relative fluxes, but the size of the emission region is the dominant factor \citep{cko, bww}. \citet{bww} have shown that measurements of the anomalous fluxes can be used to set limits on the size of the emission regions. The final method uses spectroscopic data to compare the shapes of emission lines and continuum spectra. If the broad emission line region has ordered kinematical motions, then differential magnification may change the shape of the emission lines of one image. In the case of a macro-imaged quasar, differences between the line spectra of the images can then provide a diagnostic for the kinematical motions of the broad emission lines. Such differences have already been observed in several sources (e.g. \citealt{Keeton, Eigenbrod08b, Sluse, Hutsemekers}). In this work we apply this last method to new integral field spectroscopy (IFS) of the multipy imaged quasar Q2237+0305 \citep{Huchra}. Q2237+0305 is an ideal candidate for differential microlensing experiments. The close proximity of the lensing galaxy, at $z_d=0.0394$ compared to the quasar redshift of $z_s=1.695$, yields a large projected Einstein Radius (ER) of $\sim 2\times 10^{17}h_{70}^{1/2} (M/M_\odot)^{1/2}$cm in the source plane. This ER, which characterizes the size-scale of magnification fluctuations, is significantly larger than its estimated continuum region size of $3\times 10^{16}h_{70}^{1/2} (M/M_\odot)^{1/2}$cm \citep{Witt}, and it is similar to the upper size limit of its CIII] and MgII BELR, determined from our previous observations of differential microlensing in this source (\citealt{Wayth}, hereafter W05). This indicates a high probability of differential microlensing between continuum and broad lines, and also a reasonable probability of differential microlensing {\it within} the BELR. Differential microlensing between the continuum and BELR in Q2237+0305 has been observed many times. Microlensing within the broad line itself has also been observed. Most notably, \citet{ecsma} have conducted VLT monitoring of Q2237+0305 that spans several years, and present high quality slit spectroscopy of both continuum and BELR change, comparable to the data presented in this paper. However IFS observations can provide more reliable spectra as they allow more careful deblending of the small-separation lensed images and lensing galaxy. Interpretation of the differences between the spectra of the four images requires a model for both the accretion disk and the spatial and kinematical structure of the BELR. Broadly, larger flux changes are expected for regions closer to the accretion disk. The velocity structure, projected along the line-of-sight to the quasar should discriminate between models where the higher velocities are far from the disk (due to an accelerating wind) and those where a high velocity might be found close to the disk (due to Keplerian motion around the central black hole). It is still not known which (if either) of these two models best describes the kinematics of the BELR. The current work does not attempt to construct a fully consistent model of the BELR, nor include a full photoionisation calculation. Rather, a more limited question is addressed: near the accretion disk where the continuum is emitted, is the BELR gas moving relatively quickly or slowly? Answering this question will help resolve the dominant physical process in the inner BELR. In Section 2 of this paper we present the observational data, and the methods of data reduction and spectra extraction. Section 3 describes the flux ratios for different emission lines and at different wavelengths for the continuum. In Section 4 we present simple models to describe the BELR kinematics. Section 5 describes the computation of the microlensing simulations which are convolved with the models. In Section 6, the results of comparing the model predictions with the data are presented. These are discussed in Section 7, with the final conclusions presented in Section 8.

We have presented observations of the gravitationally lensed quasar Q2237+0305 taken with the Gemini North GMOS IFU. Ratios between continuum subtracted image B and image A spectra reveal differential microlensing across the velocity structure of the CIII] broad emission line. The high velocity wings of this line tend towards the flux ratio of the continuum, and the lower velocity core, while still microlensed, is closer to the expected flux ratio in the absence of microlensing. This implies that the high velocity component is emitted from a region with a size comparable to that of the continuum emission region, whereas the low velocity component is emitted from a larger region. We conducted microlensing simulations using two simple models of the broad emission line region: an outflow model and an orbital model. The outflow model assumes a clumpy wind accellerated by radiation pressure.The orbital model assumes circular Keplerian orbits with random orientations. For both models we tested a wide range of parameters. These models were used to construct an ensemble of simulated B/A flux ratios as a function of velocity, for comparison with the observed flux ratio spectrum. A purely radial outflow was unable to reproduce the observed differential microlensing signature for any plausible launching radius of the wind. Conversely, the orbital model was able to reproduce the observed signature for all simulated black hole masses $M_{bh}< 2\times10^9M_{\odot}$. Though our orbital model is simplistic, we interpret this result as further evidence that the inner regions of the BELR are gravitationally dominated. This is consistent with an outflow model where the BELR gas is lifted off the quasar accretion disk, and thus retains a high Keplerian velocity. The BELR models presented here are not intended to accurately describe the physical situation in the quasar. They describe only the generic behaviour of a radially outflowing wind, or a collection of orbiting clouds. More sophisticated models, perhaps making use of radiative transfer codes such as CLOUDY, may provide additional constraints on the quasar central engine. We have obtained Gemini IFU observations of nine other gravitationally lensed quasars, both double and quadruply imaged. These data will allow us to probe a range of quasar orientations, emission lines, black hole masses, and therefore BELR emission region scales. The analyses of these data are forthcoming.

1012.3480

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1012.3163_arXiv.txt

We investigate the dynamical status of the low-mass globular cluster Palomar~13 by means of $N$-body computations to test whether its unusually high mass-to-light ratio of about 40 and its peculiarly shallow surface density profile can be caused by tidal shocking. Alternatively, we test -- by varying the assumed proper motion -- if the orbital phase of Palomar~13 within its orbit about the Milky Way can influence its appearance and thus may be the origin of these peculiarities, as has been suggested by \citet{Kuepper10b}. We find that, of these two scenarios, only the latter can explain the observed mass-to-light ratio and surface density profile. We note, however, that the particular orbit that best reproduces those observed parameters has a proper motion inconsistent with the available literature value. We discuss this discrepancy and suggest that it may be caused by an underestimation of the observational uncertainties in the proper motion determination. We demonstrate that Palomar ~13 is most likely near apogalacticon, which makes the cluster appear supervirial and blown-up due to orbital compression of its tidal debris. Since the satellites of the Milky Way are on average closer to apo- than perigalacticon, their internal dynamics may be influenced by the same effect, and we advocate that this needs to be taken into account when interpreting their kinematical data. Moreover, we briefly discuss the influence of a possible binary population on such measurements.

\label{Sec:Introduction} There are many objects on the sky, especially in the halo of the Milky Way (MW), whose nature is not clear to us. Some of those objects are hard to address observationally, and for others there is just no conclusive theoretical explanation. In fact, simply by looking at a colour-selected sample of stars within a region of the sky it is sometimes not easy to determine the true extent of a stellar system, mostly since it lacks a clear cut-off in its surface density profile. The same holds true for the determination of its velocity dispersion through a sub-sample of stars with readily measured radial velocities. These uncertainties typically result in discussions and speculations about a best-fitting density profile as well as a system's true tidal radius (e.g. \citealt{King66, Elson87, McLaughlin05}), and also about the true mass-to-light ratios of such systems (e.g. \citealt{Kroupa97, Mieske08}). Some of these uncertainties arise from peculiar surface density profiles. That is, even though many objects in the Milky Way halo are well limited and show a well defined surface density profile with a slope of about $R^{-4}$ in the region of the tidal radius, some objects obey shallow surface density profiles in the outskirts, having slopes of about -1 to -2, like for example the MW globular clusters Palomar 5 \citep{Odenkirchen03}, NGC 5466, M 15, M 53, M 30, and NGC 5053 \citep{Chun10}, AM 4 \citep{Carraro07}, Whiting 1 \citep{Carraro09}, and NGC 1851 \citep{Olszewski09}. The latter furthermore seems to be surrounded by a 500 pc halo of stars whose origin is unknown up to now. Other uncertainties arise from unusual mass-to-light (M/L) ratios of some stellar systems. While most globular clusters show mass-to-light ratios of 1-2, Ultra-Compact Dwarf galaxies (UCDs) have higher M/L by a factor of about two, whereas dwarf spheroidal galaxies even show values of up to $10^3$ \citep{Dabringhausen08, Geha09}. These differences are usually ascribed to different dark matter contents, catastrophic tidal heating by gravitational shocks, a variation of the IMF, tidally reshaped stellar phase-space distribution functions, contaminations from stellar streams in the MW halo, or alternative gravitational theories (e.g. \citealt{Kroupa97}, \citealt{Gilmore07}, \citealt{Simon07}, \citealt{Mieske08}, \citealt{Angus08}, \citealt{Niederste09}). The low-mass Galactic globular cluster Palomar 13 is a stellar system which shows both, an unclear extent due to a shallow surface density profile, and a high velocity dispersion resulting in a mass-to-light ratio of about 40 \citep{Siegel01, Cote02}. Further details on this cluster are presented in Sec.~\ref{Sec:Palomar 13}. In this investigation we demonstrate by means of $N$-body calculations how these observational results can be explained without the need for dark matter, tidal heating, binaries or changes in the law of gravity. To this end we compute models of Palomar 13 on various orbits about the Galaxy that are consistent with its present-day distance and radial velocity with respect to the Sun. We show how different such a stellar system can appear in different phases of its orbit. Details on the models are described in Sec.~\ref{Sec:Models}. The results of these computations and the mock observations in which we show how this cluster may appear when observed with an 8m-class telescope, are shown in Sec.~\ref{Sec:Results}. Sec.~\ref{sec:discussion} is a short discussion on the plausibility of our findings. Finally, in Sec.~\ref{Sec:Conclusions} we give a short summary and conclusions.

\label{Sec:Conclusions} We performed a set of $3\times15$ $N$-body computations of the low-mass Milky-Way globular cluster Palomar 13 which shows some peculiarities in observations. First of all, \citet{Cote02} measured a velocity dispersion of $2.2\pm0.4$ km/s, which yields a very high mass-to-light ratio of about 40 due to the cluster's low integrated absolute magnitude of only $M_V = -3.8$ mag. Secondly, Pal~13 shows a shallow slope at large radii within its surface density profile of $\eta = 1.9$, making a determination of its tidal radius difficult \citep{Siegel01, Cote02}. It has been suggested in both publications that these effects might be due to Pal~13's last pericentre passage on its eccentric orbit about the Galactic centre. In contrast, \citet{Kuepper10b} find by means of $N$-body computations that pericentre passages barely influence the appearance of a cluster's surface density profile. Instead, they find a flattening of the surface density profile only for clusters on eccentric orbits which are about to reach their apogalacticon, resulting from the compression of the tidal debris as it gets decelerated in its orbit. We therefore use three different orbits to explicitly test these two hypotheses for Pal 13 (Fig.~\ref{orbits}). First, we use the radial velocity measured by \citet{Cote02} in combination with the proper motion measured by \citet{Siegel01}. This yields an orbit in which Pal~13 today is close to its perigalacticon (\textit{orbit~1}). Secondly, we use only the measured radial velocity but set the proper motion to zero. This yields a similar orbit which is less eccentric and in which Pal~13 is today closer to apogalacticon but still not close enough for the effect of tidal debris compression described in \citet{Kuepper10b}, and which causes the surface density profile at large radii to become shallow, to take place (\textit{orbit~2}). Finally, we use the orbit with the proper motion which minimizes the orbital velocity of Pal~13, since this yields the orbit in which Pal~13 is today closest to apogalacticon (\textit{orbit~3}). As it turns out, the model clusters on \textit{orbit~3} can readily reproduce Pal~13's peculiarities both in terms of surface density profile and velocity dispersion (and thus mass-to-light ratio), whereas the two other orbits cannot (Fig.~\ref{results}). While the first two orbits yield clusters with regular equilibrium velocity dispersion, the last orbit yields an enhanced velocity dispersions and a much larger spread in velocity dispersion values when measured from a subset of 21 stars. This is due to unbound stars within the cluster (potential escapers), and extra-tidal stars which get pushed back into the vicinity of the cluster when the cluster-tail system gets decelerated on its way to apogalacticon \citep{Kuepper10a, Kuepper10b}. With this investigation we would like to stress the importance of the orbital phase of a cluster on its appearance. Particularly interesting is the orbital phase just before reaching apogalacticon, where the cluster and its tails get decelerated and thus compressed such that the stellar density, especially in the region around the cluster's tidal radius, gets enhanced with unbound stars. These stars can alter the slope of the surface density profile at large radii, and increase the measured velocity dispersion significantly. Since any cluster (or satellite in general) on an eccentric orbit about a galaxy spends most of its lifetime close to apogalacticon, it is likely to be observed in such a phase. It is therefore expected that a good fraction of all satellites are affected by this effect of orbital compression of their tidal debris. Observations not taking this effect into account may therefore assume too large tidal radii and/or ascribe a pronounced tidal debris to tidal shocking which in reality is only due to the deceleration of the satellite-tail system. And in some cases it may even lead to drastic overestimates of the dynamical mass, as is demonstrated here for Pal~13. Moreover, orbital compression of a satellite's tidal debris can produce stellar systems, not only star clusters but also dwarf galaxies, which may appear largely extended, irregular in shape and dynamically hot, and thus may be misinterpreted, for instance, as bound systems embedded in dark matter haloes. Whether this effect can explain, for instance, the high mass-to-light ratios of dwarf galaxies such as Segue 1, which is currently known as the darkest ultra-faint dwarf galaxy \citep{Geha09}, has to be checked in a future investigation focussing on typical dwarf galaxy orbits and morphologies. Since dwarf galaxies are in general more diluted than globular clusters we expect that the effect is more pronounced in those cases. Available work in this direction by \citet{Kroupa97} and \citet{Klessen98} indeed suggests similar issues are relevant for dwarf spheroidal galaxies (see also \citealt{Kroupa10}). Segue 1, in fact, has recently been found in SDSS data to show a prominent tidal debris and therefore was re-classified as a dissolving star cluster \citep{Niederste09}. Its enhanced M/L ratio was interpreted as contaminated by stars from the Sagittarius stream but may well be due to the orbital compression effect described here. Finally, this investigation poses the question how reliable proper motion measurement for halo satellites are (especially if they suffer from the orbital compression effect described in this work), or if we understand the potential of the Milky Way correctly, i.e. is the \citet{Allen91} potential which we used in this investigation a sufficient approximation? If the observed peculiarities in Pal~13 are indeed due to the orbital phase of the cluster, then either the proper motion measurement or the Galactic potential will be in question.

1012.3163

1012

1012.1957_arXiv.txt

We investigate the validity of the approximate method to describe a strong gravitational lensing which was extended by Alard on the basis of a perturbative approach to an Einstein ring. Adopting an elliptical Navarro-Frenk-White (NFW) lens model, we demonstrate how the approximate method works, focusing on the shape of the image, the magnification, caustics, and the critical line. Simplicity of the approximate method enables us to investigate the lensing phenomena in an analytic way. We derive simple approximate formulas which characterise a lens system near the Einstein ring.

Cold dark matter is one of the most important components in the universe. The cosmic microwave background anisotropies and the large scale distribution of galaxies cannot be naturally explained without the cold dark matter component. The mean density parameter of the cold dark matter has been measured precisely,\cite{WMAP} but its true character has not been identified. The elementary particle physics predicts possible candidates of the cold dark matter, and many experiments are ongoing aiming at a direct measurement. The cold dark matter is considered to be distributed associated with each galaxy, forming dark matter halo. Then, the investigation of the structure of the halos is quite important in exploring the nature and the origin of the dark matter. The strong gravitational lensing is a useful probe of the halo-structure (see, e.g.,\cite{ZR} for a review). Especially, a lens system near Einstein ring is useful because a wealth of information can be obtained.\cite{Kochanek} Besides, the strong lensing systems are also useful as a tool of the dark energy study.\cite{FutamaseII}\cdash\cite{Futamase} Because of the recent observational developments, many strong lensing systems have been found. The strong lensing statistics is now becoming one of the powerful tool for exploring the nature of the dark energy.\cite{Oguri} Future dark energy surveys will detect much more strong lensing systems (see, e.g.,\cite{LSST}), and the strong lensing system will play a more important roll in cosmology. In realistic situations, the mass distribution in a halo is not simple, which makes reconstruction of the lens model complicated. The lens equation is complicated for a non-spherical lens model, which needs to be solved numerically. Then, analytic approximate approach to strong lensing system is useful, if its validity and accuracy are guaranteed. A perturbative approach to the lensing system close to the Einstein ring configuration was developed, e.g.,\cite{Blandford,SEF}. Recently, Alard extended the perturbative approach, which is applied to analyse lensing systems.\cite{Alard07}\cdash\cite{Alardnew} In the present paper, we investigate the validity of the perturbative approach to the lensing system close to an Einstein ring, assuming an elliptical lens model. We demonstrate the validity of the perturbative approach quantitatively, by comparing with an exact approach on the basis of the numerical method, focusing on the shape of the image, the magnification, the caustics, and the critical line. Using the approximate method, expanded in terms of the ellipticity parameter of the lens model, we derive simple approximate formulas which characterise an elliptical lensing system near the Einstein ring in an analytic way. This paper is organized as follows: In section 2, we briefly review the basic formulas for the gravitational lensing and the perturbative approach to a perturbed Einstein ring, based on the work by Alard.\cite{Alard07,Alard08} In section 3, we compare the perturbative approach with the exact approach that relies on a numerical method, focusing on the shape of lensed images, the caustics, the critical curve, and the magnification, respectively. We demonstrate the validity of the perturbative approach at a quantitative level. In section 4, some useful formulas are presented, which are derived using the perturbative approach in the analytic manner. Section 5 is devoted to summary and conclusions. Throughout the paper, we use the unit in which the speed of light equals 1.

We studied the perturbative approach to the strong lensing system, which was extended by Alard, in both the analytic and numerical manners. We investigated the validity of the perturbative approach by comparing with the exact approach on the basis of the numerical method, focusing on the shape of the image, the magnification, the caustics, and the critical line. The perturbative approach works well in the case when the ellipticity of the lens potential $\eta$ is small and the configuration of the source is close to the that of an Einstein ring. At a quantitative level, the perturbative approach is valid at the $10$ percent level for $\widetilde{\delta y_{10}}\simlt0.2$ and $\eta\simlt 0.3$. We also demonstrated that the lowest-order expansion in terms of $\eta$ also works well, which enables us to investigate the lensing system in an analytic way. We investigated the critical behaviour of the lensed images, by demonstrating the phase diagram of the different configurations of four separated images (type I), an arc and one separated image (type III), and one connected ring image (type V). The critical configuration of type II appears during the transition from the type I to the type III, while the type IV appears during the transition from the type III to the type V. We investigated how the critical behaviour depends on the lens ellipticity, the source position and the source radius. We also demonstrated how the appearance of the critical configuration III and V is related to the condition between the source configuration and the caustics. The condition of the critical configuration was investigated in an analytic manner using the lowest-order expansion of the ellipticity $\eta$ of the elliptical lens potential. The perturbative approach with the lowest-order expansion with respect to $\eta$ is useful to find the simple formulas which characterises the lensing system in an analytic manner. We derived the analytical formulas of the arc width and the magnification factor. In the point source limit, the simple formulas for the image position and magnification factor were obtained. These formulas can be easily solved, which gave the simple analytic expressions in the absence from the inclination angle. These results will be useful to understand the gravitational lensing phenomena. In a realistic situation in reconstructing a gravitational lensing system, substructures in the lens might have to be taken into account. In the reference,\cite{Alard08} Alard considered how a substructure affects a lensed image in the perturbative approach. Even a substructure with small mass could make a change in the caustics and the lensed image drastically. It is an interesting problem how one can determine the gravitational lens potential including substructures simultaneously. Here, there is potentially a lot of room for improvement.\cite{Alard09} This issue is outside the scope of the present paper, but need to be elaborated for a precise reconstruction of a gravitational lens system.

1012.1957

1012

1012.1166_arXiv.txt

{Directional detection of non-baryonic Dark Matter is a promising search strategy for discriminating WIMP events from background. However, this strategy requires both a precise measurement of the energy down to a few keV and 3D reconstruction of tracks down to a few mm. To achieve this goal, the MIMAC project has been developed. It is based on a gaseous micro-TPC matrix, filled with $\rm ^3He$, $\rm CF_4$ and/or $\rm C_4H_{10}$. The first results on low energy nuclear recoils ($\rm ^1H$ and $\rm ^{19}F $) obtained with a low mono-energetic neutron field are presented. The discovery potential of this search strategy is discussed and illustrated by a realistic case accessible to MIMAC.} \FullConference{Identification of Dark Matter 2010-IDM2010\\ July 26-30, 2010\\ Montpellier France} \begin{document}

Directional detection of Dark Matter is based on the fact that the solar system moves with respect to the center of our galaxy with a mean velocity of roughly 220 km/s \cite{spergel}. Taking into account the hypothesis of the existence of a galactic halo of DM formed by WIMPs (Weakly Interacting Particles) with a negligible rotation velocity, we can expect a privileged direction for the nuclear recoils in our detector, coming out from elastic collision with those WIMPs. The MIMAC (MIcro-tpc MAtrix of Chambers) detector project tries to get these elusive events by a double detection: ionization and track, at low gas pressure with low mass target nuclei (H, 19F, 3He). In order to have a significant cross section we explore the axial, spin dependant, interaction on odd nuclei. The very weak correlation between the neutralino-nucleon scalar cross section and the axial one, as it was shown in \cite{PLB}, makes this research, at the same time, complementary to the massive target experiments.

Directional detection is a promising search strategy to discover galactic dark matter. The MIMAC detector provides the energy of a recoiling nucleus and the reconstruction of its 3D track. The first 3D tracks observed with the MIMAC prototype were shown: 5.9 keV electrons (typical background) and low energy proton and fluorine recoils (typical signal). The next step will be to build a demonstrator of 1 m$^3$ to show that the large micro-tpc matrix for directional detection of dark matter search is accessible.

1012.1166

1012

1012.1216_arXiv.txt

It is commonly believed that most of the stars born in associations decaying with characteristic velocities of stars $\sim 10$~km$/$s. For dwarf galaxies the decay can lead to ejection of stars from the galaxy. The effect is studied for spheroidal and disk dwarf galaxies, and is shown to have substantional observational consequences for disk galaxies with escape velocities up to $20$~km$/$s, or dynamical masses up to $10^8 M_\odot$. The ejection of stars can (i) reduce the abundances of the products of Type Ia supernovae and, to a lesser degree, Type II supernovae, in disk stars, (ii) chemically enrich the galactic halo and intergalactic medium, (iii) lead to the loss of $50\%$ of the stellar mass in galaxies with masses $\sim 10^7 M_\odot$ and the loss of all stars in systems with masses $10^5 M_\odot$, (iv) increase the mass-to-luminosity ratio of the galaxies.

Mass exchange between a galaxy and the intergalactic medium(IGM) can influence both the chemical composition of the galactic gas and the morphology of the galaxy. Several mechanisms for gas loss by a galaxy can be distinguished \citep{Shustov--1997A&A...317..397S}: galactic wind induced by numerous supernova explosions, ram pressure exerted by the IGM, the tidal influence of other galaxies of the group, evaporation of gas due to interactions with the hot IGM, and blowing of the gas out of the galaxy by stellar radiation. Possible ways for the loss of stellar mass include tidal interactions, the ejection of stars due to the statistical mechanism, and the decay of stellar associations. Let us consider some of these mechanisms in more detail. There have been many studies of the effects of numerous supernova explosions, such as galactic fountains, super-bubbles, and winds, on disk galaxies (see \citet{Shustov--1997A&A...317..397S,Cooper--2008ApJ...674..157C} and references therein). The efficiency of these processes for gas ejection depends strongly on the distribution of the gas. Thus, models of galaxies with a stratified interstellar medium (ISM) display a much higher mass-loss efficiency than do models with a continuous distribution of their ISM \citep{Cooper--2008ApJ...674..157C}. Computations of models with a continuous gas distribution, in turn, provide different results depending on the distribution law \citep{Mac_Low--1989ApJ...337..141M}. It was shown in the theoretical study of \citet{De_Young--1990ApJ...356L..15D} that the fraction of expelled gas in a $1.4 \times 10^9~M_\odot$ galaxy is $\sim 0.6$, but, as the authors note, the presence of dark matter was not taken into account. According to \citet{Igumenshchev--1990A&A...234..396I}, galaxies with masses exceeding $10^{12}~M_\odot$ do not have winds, and, hence, do not lose gas via this mechanism. Based on these computations, \citet{Shustov--1997A&A...317..397S} used a simple approximation for the relation between the mass of the galaxy and the fraction of expelled matter in their models: \begin{equation} \label{eq:galactic-wind} f_\mathrm{esc} = 2.4 - 0.2 \lg\frac{M_\mathrm{G}}{M_\odot} \;. \end{equation} In this approximation, the efficiency of gas ejection becomes unity for galaxies with masses of $10^7~M_\odot$, irrespective of their morphology. According to this model, galaxies with such masses should not contain gas. At the same time, gas is almost completely absent only in spheroidal and elliptical dwarf galaxies, whereas it can constitute a substantial fraction of the masses of disk and irregular galaxies \citep{Begum--2004A&A...413..525B,Karachentsev--2004AJ....127.2031K, Begum--2008MNRAS.383..809B}. On the other hand, observations have not revealed gas outflows into the IGM from galaxies with dynamical masses of $\sim 10^9~M_\odot$ \citep{van_Eymeren--2009A&A...493..511V}. Thus, the question of the efficiency of gas ejection due to galactic winds remains open. Tidal interaction may be responsible not only for mass exchange between galaxies during collisions or close fly-bys and between galaxies and the IGM, but also for changes in galactic morphology. According to the estimates of \citet{Tutukov--2006ARep...50..439T}, every galaxy in a cluster experiences a collision at least once during its lifetime. In these collisions, the galaxies may merge, lose their gaseous components, or be disrupted completely. A new galaxy may also form from gas lost in galaxy collisions. The ram pressure of the IGM gas, evaporation of gas, and sweeping-out of dust are less efficient galactic mass-loss mechanisms, though they influence the chemical evolution of galaxies and enrichment of the IGM. The essence of the statistical mechanism is that, in the case of an equilibrium distribution of the stars in the gravitational potential of the galaxy, there will always be stars with velocities exceeding the escape velocity. As these stars leave the potential well, the system relaxes to a new equilibrium state. However, the timescale for the statistical mechanism is very large --- close to a hundred relaxation times \citep{Binney--1987gady.book.....B}, where the latter is \begin{equation} \tau_\mathrm{relax} \sim \frac{0.1 N}{\ln N}\,\tau_\mathrm{dyn} \;, \end{equation} where $N$ is the number of stars in the system and $\tau_\mathrm{dyn}$ the dynamical timescale of the system. Typical galactic dynamical timescales are $\sim 10^7 \-- 10^8$~yr. Even for $\sim 10^6~M_\odot$ galaxies, the relaxation time exceeds the Hubble time. Other mass-loss mechanisms rewiewed by \citet{Binney--1987gady.book.....B} for collisionless stellar systems are even less efficient. It is commonly believed that most stars are born in associations (see however, the paper by \citet{Elmegreen--1996ApJ...466..802E}). The lifetimes of OB associations from birth to decay is short, of the order of several million years. Typical velocities of stars acquired during the decay are of the order of $10$~km$/$s, according to various studies \citep{Gvaramadze--2008A&A...490.1071G} and observations \citep{Gies--1987ApJS...64..545G}. Other estimates limit the velocity range to $2\--8$~km$/$s \citep{Brown--1997MNRAS.285..479B}. The virial velocities in low-mass galaxies can be several km$/$s \citep{Karachentsev--2004AJ....127.2031K}, and the escape velocity lower than $20$~km$/$s \citep{Bovill--2009ApJ...693.1859B,Dijkstra--2004ApJ...601..666D}. In the case of disk galaxies, the ordered motions of the galactic matter may facilitate the ejection of stars. The aim of our current study is to estimate this effect and observational manifestations in dwarf galaxies. In Section 1, we compute the probability of ejection of stars from spheroidal and disk galaxies. In Section 2, we present the results of modeling the evolution of dwarf disk galaxies taking into account the ejection of stars. Section 3 discusses our results.

We have studied the influence of the loss of stellar mass on the evolution of dwarf spheroidal and disk galaxies. The decay of OB associations was considered as a possible mass-loss mechanism, with the decay enabling some stars to obtain velocities sufficient to escape their galaxy. The decay of associations is essentially of no importance for the evolution of spheroidal galaxies. The effect is also small for disk galaxies with $\sigma_\mathrm{OB} = 2$~km$/$s. Since a value of $\sim 10$~km/s is thought to be typical, we focused our analysis on models with $\sigma_\mathrm{OB} = 8$~km/s. The results of our analysis are as follows. \begin{enumerate} \item During the lifetime of an OB association ($\sim 10^7$~yr), the most massive SN II ($\gtrsim 13~M_\odot$) are able to enrich the ISM in the products of their explosions. Lower-mass stars that leave their galaxy do not contribute to the enrichment of the disk ISM, but instead serve as a source of elements for the halo or IGM. The same is true of SNIa. \item Disk galaxies that had at the onset of their star formation masses of $3 \times 10^7~M_\odot$ contain half of their mass in disk stars and the other half in the halo. The halo luminosities in such galaxies exceed the disk luminosities by a factor $1.5 \-– 2$. We can thus infer that galaxies with masses $\lesssim 10^7~M_\odot$ that were initially disk galaxies change their morphology to spheroidal. According to the classificaiton of \citet{de_Vaucouleurs--1991trcb.book.....D}, spiral galaxies are assigned morphological indices T = 4 (see also \citet{Corwin--1994AJ....108.2128C}). In the catalog of nearby galaxies \citep{Karachentsev--2004AJ....127.2031K}, disk galaxies have absolute magnitudes not exceeding $-13^m$ (for morphological indices from 0 to 7; i.e., including lenticular and irregular galaxies that are closest to disk galaxies); this magnitude corresponds to a luminosity $\sim 10^7~L_\odot$ and a galactic mass $\sim 10^8~M_\odot$. Lower-mass galaxies are classified as spheroidal and irregular. This is confirmed by the computations for the models we have adopted here. \item In systems with masses $\lesssim 10^5~M_\odot$, a large fraction of the stellar mass leaves not only the disk of the galaxy, but also the halo (Fig. \ref{fig:mass-beta-chi}, right panel). Thus, if extremely low-mass galaxies can form at all, they can lose almost all their stellar population to the IGM after their first burst of star formation. As a result, a dark-matter halo enriched in gas should be left. This scenario may be important for the problem of missing satellites of the Galaxy. \item The ejection of stars increases the mass-toluminosity ratio. For galaxies with total masses (disk + halo) of $\sim 10^7~M_\odot$, this ratio increases by a factor of $2 \-- 2.5$ (Fig. \ref{fig:trends}, upper right panel). \item The ejection of stars may result in strong variations of elemental abundances in the gas (Fig. \ref{fig:trends},middle row of plots): along with the systematic decrease of the abundances by $\sim 0.05$ dex, the elemental abundances in the lowest-mass galaxies can increase by $0.1 \-- 0.15$ dex. The abundances in stars systematically fall with the galaxy mass decrease, by $0.2$ dex. \end{enumerate} This study was supported by the Federal Agency on Science and Innovation (state contract no. 02.740.11.0247), the Federal Education Agency (contract RNP-2.1.1-1937), the Program of State Support for Leading Scientific Schools of the Russian Federation (grant no. NSh-4354.2008.2), and the Russian Foundation for Basic Research (project nos. 08-02-91321-IND and 07-02-00454).

1012.1216

1012

1012.3213_arXiv.txt

The radiative efficiency of AGN is commonly estimated based on the total mass accreted and the total AGN light emitted per unit volume in the universe integrated over time (the Soltan argument). In individual AGN, thin accretion disk model spectral fits can be used to deduce the absolute accretion rate $\Mdot$, if the black hole mass $\Mbh$ is known. The radiative efficiency $\eta$ is then set by the ratio of the bolometric luminosity $L_{\rm bol}$ to $\Mdot c^2$. We apply this method to determine $\eta$ in a sample of 80 PG quasars with well determined $L_{\rm bol}$, where $\Mdot$ is set by thin accretion disk model fits to the optical luminosity density, and the $\Mbh$ determination based on the bulge stellar velocity dispersion (13 objects) or the broad line region (BLR). For the BLR-based masses, we derive a mean $\log \eta=-1.05\pm 0.52$ consistent with the Soltan argument based estimates. We find a strong correlation of $\eta$ with $\Mbh$, rising from $\eta\sim 0.03$ at $\Mbh=10^7\Msun$ and $L/L_{\rm Edd}\sim 1$ to $\eta\sim 0.4$ at $\Mbh=10^9\Msun$ and $L/L_{\rm Edd}\sim 0.3$. This trend is related to the overall uniformity of $L_{\rm opt}/L_{\rm bol}$ in our sample, particularly the lack of the expected increase in $L_{\rm opt}/L_{\rm bol}$ with increasing $\Mbh$ (and decreasing $L/L_{\rm Edd}$), which is a generic property of thermal disk emission at fixed $\eta$. The significant uncertainty in the $\Mbh$ determination is not large enough to remove the correlation. The rising $\eta$ with $\Mbh$ may imply a rise in the black hole spin with $\Mbh$, as proposed based on other indirect arguments.

Material falling in nearly circular orbits onto a black hole (hereafter, BH) looses a fraction of its rest mass energy during the infall. The lost energy is emitted as an outflow of radiation and particles (and potentially Poynting flux). A measurement of the fraction of mass inflow $\Mdot$ converted to radiation $L_{\rm bol}$, provides a measure of the radiation efficiency $\eta\equiv L_{\rm bol}/\Mdot c^2$. In the ``standard'' accretion disk (AD) model \citep{ss73,nt73}, the BH spin $a_*$ determines $\eta$ because it sets the marginally stable orbit, $r_{\rm ms}$, beyond which the material is assumed to fall into the BH without loosing further energy. Since the total efficiency is a rising monotonic function of $a_*$, the measured $\eta$ provides a lower limit on $a_*$. The value of $a_*$ is important because it tells us how the BH mass $\Mbh$ grew. If it grew mostly through a single event (major merger or continuous gas accretion) then $a_*$ will be close to unity. If it grew through a series of independent events (minor mergers, episodic accretions), then $a_*$ will be close to zero \citep{hb03,gsm04,vol05,kp06,bv08}. In the absence of torques near $r_{\rm ms}$, the value of $a_*$ sets, through the value of $r_{\rm ms}$, the spectral energy distribution of the accretion disk \citep{cun75,ks84,sm89,ln89,sk98,hub00}, the rotation of the polarization angle of the AD emission, \citep{csp80,lnp90,dov08,sk09}, and the profile of lines emitted by the AD \citep{fab89,koj91,lao91,dab97,bd04,br06,rf08}. These methods are currently limited by the available quality of the data, and by potential uncertainties in our models of the AD structure. As a result, we do not yet have precision measurements of $a_*$ in more than a few objects. Thus, an additional independent constraint on $a_*$ based on $\eta$, is useful. A determination of $\eta$ can potentially provide an upper limit on the additional power which may be generated by the accretion in a jet/wind outflow. Without torques, there is an upper limit on the total efficiency of 40\% for $a_*=1$, or 31\%, for the maximal spin within an AD of $a_*=0.998$ \citep{tho74}. Such outflows are important as they can couple to the surrounding gas more efficiently than radiation, and may significantly affect the host galaxy evolution (e.g. \citealt{mn07}), suppress cluster cooling flows (e.g. \citealt{chu02,all06}), and may be relevant to the correlation of the black hole mass with the galaxy properties \citep[ and citations thereafter]{mag98}. The implied jet power of AGN in cooling flow clusters can be significantly larger or smaller than the radiative power, depending on the AGN luminosity (e.g. \citealt{sha08b,mh08,cb09}). Clearly, it is useful to get an independent upper limit on the ratio of mechanical/radiative power, based on a direct determination of $\eta$. If magnetohydrodynamics torques \citep{gam99,kro99a,mg04,dhk03} are present, then the maximum efficiency can (instantaneously) exceed the limits for a no-torque disk and even exceed unity (see e.g. \citealt{ak00}) as the flow taps the spin energy of the BH. Therefore, credible estimates of such large efficiencies would provide evidence that such torques are present in real accretion flows. \citet{sol82} noted that the global AGN average radiative efficiency, $\eta_{\rm av}$, can be estimated for the AGN population by comparing the integrated $\Mbh$ per unit volume at the current epoch, with the integrated AGN luminosity per unit volume over time. \citet{sol82} also showed that $\eta_{\rm av}$ is elegantly independent of the cosmological model (a major unknown at that time). Recent studies based on the Soltan argument lead to $\eta_{\rm av}\ga 0.1$ (e.g. \citealt{yt02,erz02,mar04,bar05}). This method has also been used to estimate the time and luminosity dependence of $\eta_{\rm av}$ through more detailed modeling (e.g. \citealt{hnh06,swm09,wan09,rf09}), but the derived values are significantly uncertain. The purpose of this paper is to discuss a method to derive $\eta$ directly in individual AGN. The method assumes that the optically emitting regions of QSOs are accretion powered and radiatively efficient, thus gravitational binding energy is dissipated and radiated locally within the AD. The corresponding thin AD models were calculated to increasing levels of details, from the simple local blackbody approximation to stellar atmosphere like models where the vertical structure and the local spectrum are calculated with increasing accuracy (see \citealt{hub00} and references therein). The integrated thin AD luminosity density $L_{\nu}$ turns out to be largely set by $\Mdot$ and $\Mbh$. Thus, one can derive $\Mdot$ based on the observed $L_{\nu}$, if $\Mbh$ is known. This method has been used previously by \citet{col02} and \citet{bz03}, using simple analytic expressions, valid at long wavelengths for the emission of an Newtonian, thin, blackbody AD (e.g. \citealt{bec87}), to determine $\Mdot$ in a sample of AGN. \citet{col02} assumed a value of $\eta$ to estimate $L_{\rm bol}$ and inferred that many low $\Mbh$ AGN must be super-Eddington accretors. \citet{bz03} estimated $L_{\rm bol}$ independently for each object, which they then used to estimate $\eta$, yielding an average $\log~\eta=-1.77\pm 0.49$ in a sample of radio-quiet AGN, and $\log~\eta=-0.90\pm 0.62$ in radio-loud AGN. Observations indicate that simple thin AD model cannot reproduce the overall SED. This is due to reprocessing (IR), Comptonization in a corona (X-ray), radiative transfer effects in the inner AD, a thick AD, etc'. Our method relies on the viability of the simple thin disk approximation in the relative outer parts of the AD, which dominate the optical emission. The above effects are likely insignificant in this outer part of the AD. Thus, the redistribution of the AD radiation by these effects will not affect the measurement of $\eta$, as long as we measure the total SED, irrespective of its exact production mechanism. Here we derive $\Mdot$ based on relatively sophisticated AD models, which include relativistic effects on the disk structure and photon propagation to the observer, and solve simultaneously for the vertical structure and radiative transfer of the disk. We apply the method to the PG quasar sample \citep{sg83}, where $L_{\rm bol}$ is estimated based on high quality optical \citep{neu87}, UV \citep{bl05}, far UV \citep{sco04}, and soft X-ray \citep{blw00} observations, and $\Mbh$ is derived based on high quality spectroscopy of the H$\beta$ region by \citet{bg92}. The paper is organized as follows, in \S 2 we review the simple analytic derivation of $\Mdot$, and demonstrate that the AD $L_{\nu}$ in the optical regime is rather well determined by the local blackbody AD models, and is only slightly modified by taking into account the vertical disk structure. We also show that the optical $L_{\nu}$ is only weakly dependent on the radial disk structure, as set by $a_*$. We then derive $\Mdot$ for our sample. In \S 3 we estimate $L_{\rm bol}$, and combined with $\Mdot$ use it to compute $\eta$. We discuss the correlation, or lack thereof, of $\eta$ with parameters of interest, particularly $\Mbh$. In \S 4 we discuss various systematic effects which can affect the value of $\eta$ and the observed correlation, in particular the uncertainty in $\Mbh$, disk inclination, optical thickness of the AD emission, self-illumination, foreground extinction, and mass outflows. We summarize our conclusions in \S 5.

We have estimated the accretion rate by fitting radiatively efficient AD model SEDs (hereafter the standard model) to the optical emission in a sample of 80 PG QSOs. This method is insensitive to properties of the accretion flow and BH spacetime near the inner edge of the disk (torques at the inner radius, BH spin, advection, etc.) as long as the emission comes from large radius. We use detailed AD model SEDs which are computed from non-LTE atmospheres and include relativistic effects on photon geodesics, but find our results are qualitatively reproduced by simple, non-relativistic local blackbody relations. The derived accretion rates are nearly insensitive to spin at low black hole mass ($\Mbh \lesssim 10^8 \Msun$) and only weakly sensitive at higher mass when the fitting is done at optical frequencies. The accretion rates are more sensitive to assumed BH masses, which are estimated using broad line region virial methods as well as masses derived from the $\Mbh-\sigma$ relation for 13 sources with bulge velocity dispersion measurements. Our sample of 80 PG QSOs was chosen because they have ample broadband coverage at optical to far UV and X-ray frequencies. This allowed us to robustly estimate the bolometric luminosity modulo some uncertainty in the unobservable EUV emission. These luminosities, combined with our estimates of the accretion rate, allow us to compute the radiative efficiency for each source in our sample. We find a mean efficiency of $\log \eta = -1 \pm 0.5$, in agreement with integral constraints derived by matching local black hole mass density to the integrated quasar luminosity function (i.e. the Soltan argument). This basic agreement suggests that the standard model can provide a reasonable first-order approximation to real accretion flows, at least at radii where most of the optical emission is produced. We find a strong correlation of efficiency with BH mass (approximately $\eta \propto \Mbh^{1/2}$ in our sample) extending from $\eta \sim 0.01$ for low masses to to $\eta \lesssim 1$ at the highest masses. This relation arises because the ratio of the optical to bolometric luminosity is roughly independent of BH mass, whereas a constant efficiency thin disk model would predict a substantial increase in this ratio as mass increases. We consider three possibilities for explaining the $\eta - \Mbh$ correlation:\\ \noindent 1) The correlation is real. It could plausibly arise from a mass dependence of the BH spin driven by the differing accretion histories of black holes at different masses. Semi-analytic models of galaxy formations that attempt to model spin distributions of supermassive BHs (e.g. \citealt{lpc09,fan09}) find spin dependencies which are qualitatively, although not quantitatively, consistent with our trend of increasing efficiency with increasing BH mass. \\ \noindent 2) One or more of the observables or input parameters to the model are incorrectly estimated. Indeed, scatter in the broad line region based BH mass estimates probably contributes to some of the trend. We argue that the highest and lowest BH masses will tend to be overestimated and underestimated, respectively. This, in turn, leads to overestimates of the efficiency at the highest masses and underestimates at lowest masses. However, it would be difficult for these correlated errors to explain all of the trend, unless the true mass distribution is very narrow ($\lesssim 1$ dex), which seem highly unlikely given the large range of bolometric luminosities in our sample and the significantly larger range of $\Mbh$ derived in other studies of AGN.\\ \noindent 3) The standard (bare) AD model does not adequately approximate the dependence of real accretion flows on mass. It has well-known difficulties reproducing the observed SEDs of AGN and microlensing sizes of emission regions. A central assumption of this work is that standard model works well at larger radii where the optical emission is predominantly produced, but that some mechanism operates very near the black hole to redistribute flux from UV to X-ray frequencies. It is possible that the standard model fails in the optical emitting regions as well and that real flows naturally give rise to a relatively constant ratio of optical to bolometric luminosity which does not depend strongly on mass. We consider a number of modifications to the standard model, including irradiation. However, no single possibility we considered seemed capable of explaining the whole trend at both low and high masses, while maintaining a nearly constant $\eta$ whose value was consistent with Soltan argument.\\ The role of errors in the mass estimates could be definitively addressed by more precise mass estimates. This would either require reducing the scatter in the BLR estimates or obtaining a larger sample of sources with well-measured bulge velocity dispersions. Differentiating the degree to which the trend reflects real evolution in efficiency with mass or inapplicability of the standard models assumptions is more difficult. Microlensing models can provide independent estimates for the efficiency in lensed QSOs \citep{mor10}, although only a handful are currently available. Alternatively, relativistically broadened Fe K$\alpha$ lines could provide spin estimates from which we could infer efficiencies, but (again) precise constraints are only available for a few sources. Some light on what controls $\eta$ may be shed by the variability of $\eta$ on timescales longer than the viscous timescales, when the AD may be quasistatic. If $L_{\rm bol}\propto L_{\rm opt}^{1.5}$, then $\eta$ is constant, which is consistent with being spin driven. Until such studies are available for a significant sample of sources, this will likely remain an open question.

1012.3213

1012

1012.4119_arXiv.txt

We discuss what hampers the rate of scientific progress in our exponentially growing world. The rapid increase in technologies leaves the growth of research result metrics far behind. The reason for this lies in the education of astronomers lacking basic computer science aspects crucially important in the data intensive science era.

Present-day astronomical instruments and large surveys produce the data flow increasing exponentially in time. The CPU power required to analyse these data is also growing with the same pace following the Moore's law; the same applies to the data storage volume per price unit. However, in astronomy we do not see the exponential avalanche of scientific results produced with this computational power. This suggests the presence of \emph{a bottleneck} somewhere in the loop: \emph{if we consider the system containing three modules ``A'', ``B'', and ``C'' so that ``A'' is connected to ``C'' via ``B'', then optimizing features in module ``A'' or ``C'' will not produce a change in the performance of the system until the performance problems in module ``B'' are addressed.} Where is the true bottleneck of the scientific computing? Astronomers as many other scientists, prefer to develop their computational codes and software systems (including database solutions) themselves often having no coding skills, insufficient background in algorithms and computational science.

1012.4119

1012

1012.0998_arXiv.txt

An extended \xmm\ observation of the Seyfert 1 galaxy \ngc\ has revealed a rich absorption line spectrum indicating the presence of a photoionised outflow with a wide range of velocities and ionisation parameter. At low continuum fluxes an emission line spectrum is well defined with both narrow and broad emission components of several abundant metal ions. The absorption line velocity structure and a broad correlation of velocity with ionisation parameter are consistent with an outflow scenario where a highly ionised, high velocity wind, perhaps launched during intermittent super-Eddington accretion, runs into the interstellar medium or previous ejecta, losing much of its kinetic energy in the resultant strong shock. We explore the possibility that a quasi-constant soft X-ray component may be evidence of this post-shock cooling. This revised view of AGN outflows is consistent with multiple minor Eddington accretion episodes creating a momentum-driven feedback linking black hole and host galaxy growth.

\ngc\ is a bright, narrow line Seyfert 1 galaxy in the Ursa Major cluster, lying at a Tully-Fisher distance of 15.2 Mpc (Russell 2003), with a heliocentric velocity of 753 km s$^{-1}$ (Verheijen 2001). The rapid and large amplitude variability (Lawrence \et\ 1985,1987) was strong early evidence that the powerful X-radiation found to be a common property of Seyfert galaxies by Ariel 5 (Cooke \et\ 1976, Elvis \et\ 1978) and Uhuru (Tananbaum \et\ 1978) observations resulted from accretion onto a central supermassive black hole. Subsequent monitoring with \xte\ revealed a different aspect of \ngc, where it occasionally lapses into extended periods of low and quasi-constant X-ray emission (Lamer \et\ 2003). More recently, high resolution X-ray spectra provided by \chandra\ and \xmm\ have shown that \ngc\ also exhibits a strong ionised outflow. An early \chandra\ HETG observation resolved two X-ray absorption line systems, with outflowing velocities of photoionised gas of $\sim$600 and $\sim$2300 km s$^{-1}$, while contemporaneous HST spectra of CIV, NV and SiIV found several absorption systems with velocities from $\sim$30 km s$^{-1}$ to $\sim$650 km s$^{-1}$ (Collinge \et 2001). Of particular interest in the context of the present analysis, where we find a clear correlation of velocity and ionisation parameter, the higher velocity X-ray component had no counterpart in the UV spectra. The \chandra\ data also showed an unresolved Fe K emission line at $\sim$6.41 keV (FWHM $\leq$2800 km s$^{-1}$). Observations of \ngc\ with \xmm\ in 2001 and 2002 coincided with periods of relatively high and low X-ray flux, offering an opportunity to further explore the complexity of its X-ray spectrum. Pounds \et\ (2004), hereafter Po04, found the hard X-ray band to be dominated by reflection from cold matter, which could also explain a non-varying, narrow Fe K fluorescent line. A soft X-ray narrow emission line spectrum evident at low continuum fluxes, with observed wavelengths consistent with the \ngc\ rest frame, indicated an extended ionised emission region, while a dominant absorption line spectrum in the high flux RGS observation of 2001 showed an ionised outflow with a line of sight velocity of $\sim$500 km s$^{-1}$. An Fe K absorption line in the EPIC spectrum indicated the presence of a more highly ionised outflow component with a velocity some 10 or 30 times higher, depending on the ionisation state being FeXXVI or FeXXV. Most recently, Steenbrugge \et\ (2009) have reported the analysis of \chandra\ LETG observations in 2001 and 2003, finding evidence for a more complex ionised outflow, with low ionisation components (log$\xi$$\sim$0.1 and log$\xi$$\sim$0.5-0.9) outflowing at v$\sim$200-330 km s$^{-1}$, a more highly ionised component of log$\xi$$\sim$2 at v$\sim$600 km s$^{-1}$ and a high ionisation component (log$\xi$$\sim$3) at v$\sim$4600 km s$^{-1}$. Those authors find no evidence for recombination in 3 of the 4 ionisation components up to 20 ks after a sudden, factor-of-5 drop in X-ray flux, indicating their location at a radial distance $\ga$$7\times 10^{16}$ cm for the higher velocity gas and $\ga$$9\times 10^{17}$ cm for the lowest velocity absorption in OV. We note that these minimum radii conflict with those from a similar time variability analysis of Krongold \et\ (2007), and have the incidental bonus of removing the conceptual problem in Krongold \et\ where the local escape velocity exceeded the measured outflow value. Evidence for much higher velocity X-ray outflows (Chartas \et\ 2002, Pounds \et\ 2003, 2006; Reeves \et\ 2003, Cappi 2006, Tombesi \et\ 2010) has been confined to the very highly ionised matter (log$\xi$$\sim$3.5--4) most readily detected in the Fe K band. The high velocities and high column densities in these Fe K band observations are thought to offer the best prospect that such energetic flows represent a significant feedback mechanism to constrain the continued growth of a black hole and star formation in the host galaxy, although to date only in the case of the bright QSO PG1211+143 has the requisite wide angle flow and large covering factor been directly observed (Pounds and Reeves 2009). In the case of \ngc\,in contrast, Krongold \et\ (2007) concluded that the low velocity gas has a small covering factor and hence involves a relatively insignificant mass and energy rate. While Steenbrugge \et\ (2009) do not confirm the high density and small radial distance of the ionised outflow, they also derive a low mass and energy rate from the soft X-ray spectrum of \ngc. In this paper we report an analysis of a new \xmm\ study of \ngc\ with substantially greater sensitivity than hitherto, finding a rich absorption line spectrum revealing a velocity-structured outflow covering a wide range in velocity and ionisation parameter. We re-examine the suggestion in Po04 that both high and low velocity absorption lines arise at different stages of a mass-conserved outflow. Our new assessment considers the new data in the context of a high velocity, highly ionised wind being slowed by interaction with the interstellar medium, losing much of its mechanical energy in the resulting shock, while retaining momentum to push low ionisation gas out into the host galaxy (King 2010). We note that the same ideas would apply to an outflow colliding with earlier, slower moving ejecta. The wider importance of such shocked outflows, perhaps launched during an Eddington accretion episode (King and Pounds 2003), lies in the possibility that the accumulated thrust from multiple episodes - rather than the outflow energy - would eventually drive gas from the bulge, thereby limiting further star formation and black hole growth. Such a momentum-driven feedback mechanism has been shown by King (2003,2005) to reproduce the observed correlation of black hole and galaxy masses (e.g. Ferrarese and Merritt 2000, Gebhardt \et\ 2000, Haring and Rix 2004)).

A striking feature of the absorption spectrum in \ngc\ is the very wide range of velocities and ion stages observed. A second feature is the correlation of velocity and ionisation parameter in the continuum absorption. Figure 9 visualises that correlation, with contrasting velocity profiles of the opacity in OVIII and OVII, for the high flux revs 1722 and 1724. The upper plot, centred at the rest wavelength of OVIII Lyman-$\alpha$, shows the onset of absorption at $\sim$7500 km s$^{-1}$, with the opacity increasing to $\sim$6000--5000 km s$^{-1}$, whereafter it decreases again to disappear at $\sim$3500 km s$^{-1}$, re-emerging strongly below $\sim$1500 km s$^{-1}$. The corresponding velocity profile in the OVII 1s-2p resonance line shows only weak high velocity opacity at $\sim$4500--3000 km s$^{-1}$. In contrast the low velocity absorption is notably stronger in the lower energy ion Figure 8 brings together the velocity data from Gaussian fitting to the RGS and EPIC absorption spectra, plotting each well-defined velocity against the optimum ionisation parameter for the parent ion. The strongest absorption components at each velocity generally lie to the right hand - higher ionisation - side of the plot, consistent with a broad correlation of velocity and ionisation parameter, while the highest velocities are only detected in the higher energy parent ions, and the low velocity group are seen most strongly - or only - in low ionisation matter. The correlation of velocity and ionisation parameter is further illustrated by the results from XSTAR modelling of the highflux absorption spectrum, represented by asterisks in figure 8, with components 1, 2 and 3 following a clear linear trend in velocity and ionisation parameter . Figure 8 also indicates that part of the low velocity absorption does not follow the general trend with ionisation, with XSTAR component 4 supporting the evidence from Gaussian fitting that low velocity absorption exists at all ionisation stages covered by the RGS data. Figure 10 shows schematically how separate absorption spectra are associated with the direct continuum and with self absorption in the post shcck gas. More detailed consideration of a low velocity absorption structure, visible across all ion stages in the RGS data (and represented by components 3, 4 and 5 in the XSTAR modelling), is deferred to Paper II where it is interpreted as self-absorption in a limb-brightened shell. Meanwhile, an overview of figure 8 suggests 3 broad velocity regimes in the outflow in line-of-sight to \ngc, including the higher velocity component in the Fe K absorption. On this overview, the individual velocity components picked out by Gaussian fitting at $\sim$4000 km s$^{-1}$, $\sim$6000 km s$^{-1}$ and $\sim$8500 km s$^{-1}$, but appearing more like a broad trough in the strongest absorption, in OVIII Lyman-$\alpha$ (figure 9), may represent density variations within a continuous flow. Further evidence for such density variations, which might represent a residual shell-structure linked to the intermittent nature of the initial fast wind, can be seen in velocity profiles taken at different flux levels for different ions. \begin{figure} \centering \includegraphics[width=6.1cm, angle=270]{graph.ps} \caption {Outflow velocities derived from the Gaussian fitting plotted against the optimum ionisation parameter for each parent ion stage. Also shown by asterisks are the parameters of the 4 photoionised absorbers derived from XSTAR modelling of the RGS absorption spectra, together with a velocity/ high ionisation point to represent the putative pre-shock wind} \end{figure} In the remainder of this paper we refer to the low ($\leq$1000 km s$^{-1}$), intermediate ($\sim$3000-9000 km s$^{-1}$) and high velocity ($\sim$30000 km s$^{-1}$) regions indicated in figure 8 and explore the possibility that they represent different stages in a shocked outflow. The implications of such an interpretation are doubly significant. The concept of a slowing and cooling/recombining ionised outflow is in contrast to most current ideas for the radial acceleration of AGN winds, while efficient post-shock cooling would mean that the mechanical energy in a fast outflow may not be the primary mechanism for AGN feedback. \begin{figure} \centering \includegraphics[width=6.1cm, angle=270]{newfig9.ps} \centering \includegraphics[width=6.1cm, angle=270]{newfig9a.ps} \caption {Contrasting velocity profiles of the opacity in OVIII and OVII, for the high flux revs 1722 and 1724, illustrate the strong correlation of velocity and ionisation parameter consistent with a decellerating post-shock flow. The upper plot, centred at the rest wavelength of OVIII Lyman-$\alpha$, shows the onset of absorption at $\sim$7500 km s$^{-1}$, with opacity increasing to $\sim$6000--5000 km s$^{-1}$, whereafter it decreases again to disappear at $\sim$3500 km s$^{-1}$, re-emerging strongly below $\sim$1500 km s$^{-1}$. The corresponding velocity profile in the OVII 1s-2p resonance line shows significant high velocity opacity only at $\sim$4500--3000 km s$^{-1}$. In contrast the low velocity absorption is notably stronger in the lower energy ion} \end{figure} \subsection{Outflows and feedback in AGN} Comparison with previous observations of \ngc\ by \xmm\ and \chandra\ suggest that the low velocity absorption is persistent, while that at higher velocities is probably variable, at least on a timescale of years. The caveat on variability is that most of the earlier observations were of lower sensitivity, while the relatively long 2001 \xmm\ observation does - on a re-examination - show evidence of blue shifted absorption corresponding to an outflow in the range $\sim$4000-6000 km s$^{-1}$. The new \xmm\ spectra are most closely approached by the \chandra\ LETG spectra reported by Steenbrugge (2009), who find outflow velocities of $\sim$200, $\sim$600 and $\sim$4600 km s$^{-1}$, modelled by 4 ionised absorbers with ionisation parameters ranging from log$\xi$$\sim$0--3. In discussing the 2001 RGS spectrum of \ngc\ Po04 noted that, with the simple assumption of conservation of mass in a radial outflow, an extended region of slow moving, low ionisation gas might be a continuation of the high velocity, high ionisation flow seen in absorption in the Fe K band. On this picture much of the mechanical energy in the initial outflow would have been lost before reaching the lower ionisation stage, and Po04 speculated that this might be due to internal shocks occurring in the high velocity gas. We now take up that idea again but in the context of a shock interaction with slower moving matter, either the local ISM or previous ejecta. To justify that approach we note that the broad correlation of velocity and ionisation state in the absorption spectra for \ngc\ is a clear signature of a cooling, decellerating and recombining outflow. The mass rate in a radial outflow of velocity v, and particle density n at radius r is $\mo = 4\pi bnr^2.v.m_{p}$, where b is the fractional collimation angle and $m_{p}$ is the proton mass. Mass conservation in the flow requires the product n.r$^{2}$.v to be constant. As n.r$^{2}$ = L$_{ion}$/$\xi$, if the ionising radiation is unchanged (or changes very little), for example over a distance small compared with r, then conservation of mass in a radial outflow yields a linear correlation of velocity and ionisation parameter. We would expect that to be the case in a post-shock cooling shell. The important implication of a decellerating radial outflow is that the mechanical energy in the flow would be substantially reduced as the flow is slowed. King (2010) has recently examined a relevant scenario, where a highly ionised, high velocity wind drives into the interstellar medium, losing much of its energy by efficient cooling of the shocked gas. Such a scenario would have major implications for studies of AGN feedback based on X-ray absorption spectra. Until now, X-ray observations have been used in attempts to show that fast, ionised outflows can provide the link between the growth of a SMBH and its host galaxy, by the integrated mechanical energy in the fast flow (eg Pounds and Reeves 2009). However, as recently pointed out (eg King 2010), if such an energetic wind persists while the black hole doubles its mass (the Salpeter time), the coupling of wind energy to galactic baryons may have to be inefficient to allow massive bulges to grow to the values observed, effective feedback instead being enabled by the total momentum of the flow, an alternative that has been shown to yield the observed M - $\sigma$ relationship for nearby active galaxies (King 2003, 2005). In the following section the new \xmm\ observations of \ngc\ are examined in the context of a shocked wind, where the intermediate velocity/intermediate ionisation outflow corresponds to the immediate post-shock gas and the low velocity/low ionisation absorption to matter building up ahead of the contact discontinuity. \subsection{Comparing the \ngc\ data with a shocked wind model} In the King (2010) shocked wind model a high velocity ionised outflow collides with the ISM of the host galaxy, resulting in a strong shock. The gas density increases by a factor $\sim 4$ at the shock front, and the velocity drops by the same factor. Beyond this (reverse, adiabatic) shock, the flow is further compressed in a relatively thin, cooling region, while the velocity slows to low values. Strong Compton cooling by the AGN radiation implies a fairly rapid transition between the immediate post-shock regime and the much slower and compressed state near the contact discontinuity. Beyond the contact discontinuity a further low velocity, low ionisation component will build up as the interstellar medium is swept up by an outer (forward) shock. Comparison of that scenario with the new \ngc\ data assumes an intermittent highly ionised wind, with typical values of v$\sim$0.1c and log$\xi$$\sim$4 (Tombesi \et\ 2010), has collided with the ISM or previous slower moving ejecta, with the density increase rendering the immediate post-shock gas, v$\sim$0.025c ($\sim$7500 km s$^{-1}$), visible as resonance absorption at log$\xi$$\leq$3. The shocked wind cools and slows after passing through the inner shock, yielding the broad absorption trough seen in the velocity plot for OVIII Lyman-$\alpha$ line (figure 9). We speculate that the observed gap between the intermediate and low velocity absorbing material arises from the falling column density of OVIII, re-emerging at low velocities as the column density builds up ahead of the contact discontinuity. It is interesting to consider how the observations of \ngc\ may be used to estimate the parameters of an earlier, perhaps intermittent, Eddington episode whose effects we may now be observing. While direct evidence for a high velocity wind is marginal in the 2009 data, as it was in the 2001 \xmm\ observation, Tombesi \et\ report the detection of an outflow at $\sim$0.13c in the 2002 observation. Taken together, those reports indicate the present fast outflow in \ngc\ is intermittent, and we note that with the unusually high column density required to detect a blue-shifted Fe K absorption line in a low redshift source such as \ngc, radial expansion would rapidly render a transient outflow undetectable with EPIC. In what follows, we explore the effects of such an intermittent fast outflow as it interacts with the ISM or slower moving ejecta, as an explanation of the velocity- and ionisation-structured outflow phases observed in the present observation of \ngc. \subsection{The intermediate ionisation/intermediate velocity gas} We identify the immediate post-shock outflow in \ngc\ with the onset of substantial opacity at a velocity v$\sim$$7\times10^{8}$ cm s$^{-1}$ (figure 9, top panel). XSTAR modelling (Table 3) finds an ionisation parameter log$\xi$$\sim$3 and column density of N$_{H}$$\sim$$1.4\times 10^{22}$ cm$^{-2}$ to represent this flow phase. Both velocity and ionisation parameter are consistent with the factor $\sim$4 change from the putative 0.1c wind expected across a strong shock. The incident ionising luminosity of $\sim$$8\times10^{41}$ erg s$^{-1}$, for an average-flux EPIC spectrum (rev 1729), together with the fitted ionisation parameter and measured velocity, give nr$^{2}$v $\sim$$6\times 10^{47}$ s$^{-1}$ for the post-shock flow. Assuming an angular collimation b=0.3, the post-shock flow mass rate is then $\mo$ $\sim$$4\times 10^{24}$ gm s$^{-1}$ ($\sim$0.06 $\msun$ yr$^{-1}$). Interestingly, that rate is close to the Eddington accretion rate for \ngc\ assuming an accretion efficiency of 0.1. The corresponding momentum rate of the intermediate velocity gas is then $\sim$$3\times 10^{33}$ (cgs), with mechanical energy (0.5$\mo.v^{2}$) of $\sim$$10^{42}$ erg s$^{-1}$. That mechanical energy rate is $\sim$0.3$\%$ of the Eddington luminosity for a black hole mass of $\sim$$1.7\times 10^{6}$\Msun (Denney \et\ 2009), a factor $\sim$30 less than the value v/c.L$_{Edd}$ predicted by a simple continuum driving model (King and Pounds 2003), roughly consistent with the expected velocity-linked loss at the strong shock. As the post-shock gas cools we can identify second and third XSTAR components on the velocity/ionisation plot (figure 8), with v$\sim$3850 km s$^{-1}$ and log$\xi$$\sim$2.55, and v$\sim$550 km s$^{-1}$ and log$\xi$$\sim$1.56, following reasonably closely the expectation of a linear correlation of velocity and ionisation parameter in a post-shock flow. The detection of several strong RRC and broad resonance line emission is indicative of a strongly recombining stage in the flow, and the observed OVII RRC flux of $\sim$$4\times 10^{-5}$ photons cm$^{-2}$s$^{-1}$ provides a measure of the intermediate velocity flow. A temperature of $\sim$5 eV, from the width of the RRC, implies a recombination rate for OVIII of $\sim$$10^{-11}$~cm$^{3}$~s$^{-1}$ (Verner and Ferland 1996). With that value and assuming 30 percent of recombinations from the majority OVIII ion direct to the ground state, we deduce an emission measure from the OVII RRC flux of order $\sim$$2\times10^{63}$cm$^{-3}$, for a Tully-Fisher distance of 15.2 Mpc. We note, in passing, that this could be a minimum measure as any much higher temperature RRC component would be difficult to resolve. Importantly, the particle density in the intermediate flow region is constrained by evidence of a narrowing of the broad emission component in OVIII in rev 1739, indicating a change in the ionisation state following a 3-4 day interval of unusually low continuum flux level. This is discussed in more detail in PaperII. In the following estimates of the flow properties we assume a recombination timescale for OVIII of $\sim$4 days, corresponding to a particle density n$\sim$$5\times 10^{5}$cm$^{-3}$ for the intermediate velocity flow. The OVII RRC emission measure then corresponds to an emitting volume of $\sim$$8\times10^{51}$cm$^{3}$ and - taking the column density of component 2 in the XSTAR modelling as a measure of the intermediate velocity absorber - gives a radial thickness of order $\Delta$r$\sim$$4\times10^{15}$ cm. At a mid-phase velocity of $\sim$4000 km s$^{-1}$, the shocked gas would traverse this cooling region in a time t$\sim$$10^{7}$ s. That would suggest the intermediate velocity absorption could exist in the absence of a fast outflow for a similar timescale, while the considerably shorter recombination timescale could explain structure in the intermediate velocity flow, in turn perhaps reflecting the intermittent nature of the fast wind. From the above estimates of particle density, shell thickness and emission volume (acknowledging that these are only a crude measure across a likely strong radial gradient), we find a shell radius r$\sim$$7\times10^{17}$ cm, for b=0.3. We now compare the above parameters of the intermediate velocity/ionisation post-shock gas with values for the low ionisation/low velocity gas accumulating ahead of the contact discontinuity, to provide an order-of-magnitude estimate of the duration and history of the putative Eddington accretion episode in \ngc. \begin{figure} \centering \includegraphics[width=7.7cm, angle=0]{shell.eps} \caption {Sketch showing the origin of separate absorption spectra, in the continuum by line-of-sight to the AGN and by self-absorption in the soft X-ray emission from a limb brightened shell} \end{figure} \subsection{Constraining the low ionisation/low velocity gas} We take component 3 from the XSTAR photoionisation modelling to represent the low ionisation/low velocity gas accumulating ahead of the contact discontinuity (CD) and seen in continuum absorption in the highflux spectrum. That component has an ionisation parameter of log$\xi$$\sim$1.43, absorbing column density of $\sim$$10^{20}$cm$^{-2}$, and outflow velocity of $\sim$530 km s$^{-1}$. With a mean ionising luminosity ($\ga$l keV) of $8\times 10^{41}$ erg s$^{-1}$, we obtain n.r$^{2}$$\sim$$3\times 10^{40}$ for the low velocity/low ionisation flow gas. With the particle density scaling from the velocity difference - and reflected in the lower ionisation parameter, we assume n$\sim$$5\times 10^{6}$cm$^{-3}$, giving a radial depth of the pre-CD shell $\delta$r $\sim$ $2\times 10^{13}$ cm$^{-2}$. The higher assumed density of the pre-CD shell would appear to conflict with the persistence of the low velocity absorption, particular in a higher level line such as OVIII Lyman-$alpha$. However, as we report in Paper II, self-absorption of the broad emission line can fully account for the low velocity opacity when the continuum is weak. As the radius of this low ionisation shell must exceed the estimated thickness of the intermediate flow region, we note that light travel time delays would further limit variability in the low velocity opacity, consistent with the lower limit of r$\ga$$9\times10^{17}$ cm obtained by Steenbrugge \et\ (2009) from a lack of variability in the low velocity component in the \chandra\ absorption spectrum. We assume r$\sim$$10^{18}$ cm below. With that geometry the mass of accumulated low ionisation gas, seen in absorption, is $\sim$$7\times 10^{32}$ gm ($\sim$0.3\Msun). Comparison with the mass rate in the post shock flow indicates an accumulation time of $\sim$6 years. For a velocity of 530 km s$^{-1}$, the mechanical energy in this component of the low velocity gas is then $\sim$$10^{48}$ergs, indicating that $\sim$99.5\% of the mechanical energy has been lost in the post-shock cooling. While the total flow momentum will be conserved through the shock, we expect a major fraction of the initial ram pressure is converted to gas pressure at the contact discontinuity. Integrating the immediate post-shock outflow momentum over 6 years totals $\sim$$10^{41}$ (cgs). In comparison the accumulated ram pressure in the post-shock, low ionisation gas is $\sim$$3.7\times10^{40}$ (cgs), indicating $\sim$63\% of the immediate post-shock momentum has been translated into pressure ahead of the contact discontinuity. \subsection{Evidence for radiation from the shocked gas} If the shocked outflow scenario does apply to \ngc\ the question arises as to whether there is direct evidence for radiation from the post-shock cooling? The high temperatures in the shock require strong cooling which is likely to be dominated by Compton scattering of the AGN's radiation field. King (2010) finds that this typically will have a Compton temperature $T_c\sim 10^7$~K, compared with the much higher adiabatic shock temperature of $m_pv^2/k$ $\sim$$10^{11}$~K. In assessing whether there is any direct evidence of such cooling radiation, we recall that Uttley \et\ (2003) identified a quasi-constant soft X-ray component, modelled as a $\Gamma$$\sim$3 power law, in a \chandra\ TOO observation made during a 6-week low flux state of \ngc. The overall X-ray spectrum in the present low flux rev1739 data is very similar to the \chandra\ observation, and also to that during the \xmm\ observation in 2002 which followed a 20-day low flux period. Furthermore, regular monitoring with Swift during the present \xmm\ campaign indicated that rev1739 also followed a low flux state lasting for several days. We therefore take the rev1739 spectrum of \ngc\ to be typical of a possible base level spectrum. Figure 11 shows the pn data from rev1739 fitted with two continuum components. Above the sharp spectral break at $\sim$1-2 keV the spectrum is parameterised by a hard power law ($\Gamma$$\sim$0.8), while the steep soft X-ray component can be modelled by either a much softer power law ($\Gamma$$\sim$3.5), or a Comptonised spectrum (as shown in the figure) with kT$\sim$0.3 keV and optical depth $\tau$$\sim$0.3. A strong Gaussian emission line (equivalent width 250$\pm$30 eV) at 6.38$\pm$0.01 keV is consistent with Fe K fluorescence from the hard spectral component being reflection dominated, as reported in Po04. The short-term flux variability seen throughout the rest of the 2009 \xmm\ observation of \ngc\ is notably absent in rev 1739 (Vaughan \et\ 2010), supporting the view that the soft X-ray component in figure 11 does indeed comprise a quasi-constant low flux emission, with the lack of variability indicating its origin in an extended region. The post-shock cooling shell could fit that description and it is notable that the present RGS spectra show that absorption is limited to self-absorption in the broad emission lines at the lowest continuum fluxes. The luminosity in the Compton component of figure 11 ($\ga$0.2 keV) is $\sim$$3\times10^{41}$ erg s$^{-1}$, comparable to the mechanical energy lost in traversing the post-shock region. We have suggested that the broad emission lines and recombination continua (RRC) observed in the RGS spectrum arise from additional 2-body cooling of the post-shock outflow. Fitting in XSPEC finds fluxes of $\sim$$1.5\times10^{-4}$ photons cm$^{-1}$ s$^{-1}$ and $\sim$$2\times10^{-5}$ photons cm$^{-1}$ s$^{-1}$, respectively, for the broad emission lines of OVII and OVIII Lyman-$\alpha$. Addition of the corresponding broad emission lines of Ne, N and C yields a total soft X-ray flux of $\ga$$10^{-3}$ photons cm$^{-1}$ s$^{-1}$, and a luminosity $\ga$$2\times10^{40}$ erg s$^{-1}$. Strong RRC of O, N and C and emission from Fe-L lines increase the total observed recombination cooling to $\sim$$5\times10^{40}$ erg s$^{-1}$, making a significant addition to the cooling in the later stages of the post-shock flow. We conclude that identifying a quasi-constant soft continuum and soft X-ray emission features in \ngc\ with the cooling of the shock outflow is reasonable on energetic grounds. \begin{figure} \centering \includegraphics[width=6.1cm, angle=270]{compt.ps} \caption {Parametric model fit to the lowflux rev1739 pn data. The hard power law and soft Comptonised continuum components are shown as dotted and dashed lines} \end{figure} \subsection{Relating \ngc\ to other AGN with powerful ionised winds} Evidence has been growing for ultra-fast winds of highly ionised matter in a number of AGN, carrying mechanical energy up to 10$\%$ of L$_{bol}$ ( Pounds and Reeves 2009). The direct determination of a large-angle flow in the QSO PG1211+143 was important in confirming that the v$\sim$0.13c wind was energetically significant in terms of AGN feedback. Indirect support for such fast outflows to be typically of wide angle has recently been obtained in a survey of bright AGN by Tombesi \et\ (2010), who find some 30\% of their sample show evidence for an ionised wind of v$\sim$0.1c. As noted above, the potential importance of such powerful winds lies in providing a feedback mechanism linking the growth (and termination of growth) of supermassive black holes in AGN with that of their host galaxy. This remains true if the impact is actually delivered by a momentum-driven thrust, as would be the case if the initial ionised wind lost much of its mechanical energy in shocks before reaching the star-forming region. The analysis outlined above suggests that may be the case for the bright, nearby Seyfert galaxy \ngc. The question of how a highly ionised gas is accelerated to such high velocities is clearly important. Continuum driving for AGN accreting at a super-Eddington rate was proposed by King and Pounds (2003), and appears readily applicable to PG1211+143. However, in the case of many of the Tombesi \et\ sample, including \ngc, current mass estimates suggest they are accreting more typically at only $\sim$10-20 \% of L$_{Edd}$. If the Tombesi \et\ findings are confirmed, with highly ionised outflows at v$\sim$0.1c being relatively common in bright nearby galaxies, then perhaps Eddington or mildly super-Eddington accretion is also more common than generally believed. A case for AGN black hole masses being over-estimated has recently been argued by King (2010a). Furthermore, if the ejection of fast outflows is intermittent, as suggested here for \ngc, then only where such a wind is current or was launched very recently will it retain a line-of-sight column close to the theoretical value for continuum driving of N$_{H}$$\sim$$10^{24}$ cm$^{-2}$. \subsection{Is the forbidden line emission dominated by swept up ISM?} Strong and narrow emission lines in the RGS data, of similar strength in both low and high flux spectra, arise from the `forbidden' transitions in the 1s-2p triplets of He-like NVI, OVII and NeIX. In each case we find the line is at best only marginally resolved, with the strongest OVII line indicating a FWHM$\leq$250 km s$^{-1}$. When adjusted for the known redshift of \ngc\ the OVII forbidden line also has a very low outflow velocity of -125$\pm$40 km s$^{-1}$. We note that the velocity width is consistent with the [OIII] line width in the NLR, of 210-330 km s$^{-1}$ FWHM (De Robertis and Osterbrock 1984). In P04 constraints on the low ionisation/low velocity gas were obtained from the OVII emission line flux by noting that the 2002 November \xmm\ observation took place some 20 days after the source entered an extended low flux state, while the emission line strength of the OVII forbidden line was essentially the same as when \ngc\ was much brighter in 2001 May. This was taken to indicate a recombination time $\ga$$2\times10^{6}$s, and a plasma density n$\leq$$8\times 10^{4}$cm$^{-3}$. We now note the forbidden line fluxes in the 2009 RGS data are consistent with those measured 7 years earlier. If that is a real measure of lack of variability it would indicate a still lower density n$\leq$$10^{3}$cm$^{-3}$, or a physical extent of the forbidden line emission region larger than the present estimate for a post shock region of radius $\sim$0.3 pc. The question then arises, does a substantial fraction of the forbidden line emission come from the slower moving swept-up ISM ahead of the outer shock? This possibility is raised by noting that the resonance absorption line of OVII, identified in the context of a shocked wind with matter ahead at the contact discontinuity, has a lowest velocity component of 440$\pm$60 km s$^{-1}$, well separated from that of the observed OVII forbidden line. Absorption in OV and - less unambiguously - in OIV also have lower velocities and lie well to the low ionisation side of the linear correlation with velocity that fits the main outflow stages in figure 8. Estimating the baryon mass of the swept-up ISM from the present data depends both on the assumed density and also - quite strongly - on the relevant ionisation parameter. While the optimum ionisation parameter for OVII is log$\xi$$\sim$0.9, the very low velocity would associate it more with absorption from OIV-VI and log$\xi$$\sim$0, where OVII would represent only $\sim$3 \% of total oxygen. With the measured OVII forbidden line flux of $\sim$$1.5\times10^{-4}$ photons cm$^{-1}$ s$^{-1}$, a distance to \ngc\ of 15.2 Mpc, 3\% of oxygen in the form of OVII, and a recombination rate at kT $\sim$3 eV of $2\times 10^{-11}$~cm$^{3}$~s$^{-1}$ (Verner and Ferland 1996) we find an emission measure of $\sim$$10^{64}$cm$^{-3}$. For a particle density n$\sim$$10^{3}$cm$^{-3}$ that emission measure would correspond to a total swept-up mass of $\sim$7500 \Msun. Assuming an ISM density close to the virial equilibrium value (i.e. the isothermal sphere value for $\sigma$ = 88 km $s^{-1}$ in \ngc), the swept up mass within r$\sim$$10^{18}$ cm, for a gas fraction of 0.15, would be $\sim$$3.5\times 10^{38}$gm ($1.7\times 10^{5}$ \Msun. The implication is that previous Eddington episodes in \ngc\ have substantially reduced the gas density in the inner core of the galaxy.

1012.0998

1012

1012.1335_arXiv.txt

We derive a general expression for the large-scale halo bias, in theories with a scale-dependent linear growth, using the excursion set formalism. Such theories include modified gravity models, and models in which the dark energy clustering is non-negligible. A scale dependence is imprinted in both the formation and evolved biases by the scale-dependent growth. Mergers are accounted for in our derivation, which thus extends earlier work which focused on passive evolution. There is a simple analytic form for the bias for those theories in which the nonlinear collapse of perturbations is approximately the same as in general relativity. As an illustration, we apply our results to a simple Yukawa modification of gravity, and use SDSS measurements of the clustering of luminous red galaxies to constrain the theory's parameters.

Introduction} The cosmological constant ($\Lambda$) + cold dark matter (CDM) + general relativity (GR) model has been very successful in accounting for current cosmological data. It is incumbent upon us to put this standard model to further tests. One of its characteristic predictions is a scale-independent sub-Hubble linear growth, which can be violated if dark energy clusters\footnote{Dark energy other than the cosmological constant generally clusters, though the degree of sub-Hubble clustering is often negligible. Exceptions include models where the dark energy sound speed is substantially less than unity, e.g.~\cite{Creminelli2009}. }, or if gravity is modified \cite{Dvali:2000hr,Nicolis2009,deRham2010,Fierz1956,Jordan1959,Brans1961,Buchdahl:1983zz,Carroll2004}. Both types of models introduce new scales into the growth of structure: the Jeans scale in the case of clustered dark energy, and the GR-to-non-GR transition scale in the case of modified gravity. The most direct way to test for this effect is to measure the matter power spectrum at different redshifts, and reconstruct the growth factor as a function of scale. Here, we focus on a corollary of a scale-dependent growth, a scale-dependent halo bias, which in principle allows us to discern scale dependence even with measurements of the large-scale structure at a single redshift. The fact that a scale dependence in growth implies scale dependence in halo bias (on linear scales) was pointed out by \cite{Hui:2007zh} (henceforth HP), generalizing earlier work by \cite{Fry:1996fg,Tegmark1998}: \begin{eqnarray} \label{HPresult} b_1 (k_0, z_{\rm obs}; z_{\rm form}) = 1 + [b_1 (z_{\rm form}) - 1] {D (k_0, z_{\rm form}) \over D (k_0, z_{\rm obs})} \end{eqnarray} where $b_1(k_0, z_{\rm obs}; z_{\rm form})$ signifies the linear bias on scale $k_0$ observed at redshift $z_{\rm obs}$, for haloes that form at redshift $z_{\rm form}$. The symbol $D$ denotes the linear growth factor at the relevant scale and redshift. This expression, which assumes passive evolution, i.e. halo number conservation after formation, tells us that the observed bias would inherit a scale dependence from the growth, even if the formation bias is scale-independent. In this paper, we wish to relax these two assumptions: scale-independent formation bias, and no mergers. The paper is organized as follows. The extended Press-Schechter, or excursion set, formalism \cite{Press:1973iz, Bardeen:1986, Bond:1990iw, Lacey:1993, Sheth:1998ew} is described and generalized to allow for a scale-dependent growth in Sec.~\ref{sec:exset}. This is used to compute the halo mass function and the halo bias. In Sec.~\ref{sec:illustration}, we describe an illustrative example, a modified gravity model of the Yukawa type, and present a calculation of the linear growth factor and the excursion barrier (collapse threshold). We present the halo bias for this example, and compare it with observations. We conclude in Sec.~\ref{sec:conclusions}. In this paper, for the purpose of illustration, we have chosen to focus on one particular model of modified gravity. Our results in \S \ref{sec:exset} for the scale-dependent halo bias are, on the other hand, fairly general. A recipe for using these results in more general settings is summarized in Sec. \ref{sec:conclusions}. Readers who are interested primarily in applications, and not on the derivation, can skip directly to Sec. \ref{sec:conclusions}. In Appendix~\ref{sec:extrap} we discuss some details concerning the use of the excursion set method for theories with a scale-dependent growth factor. In Appendix \ref{sec:scalartensor} we give the derivation of the Yukawa model from a scalar-tensor theory, and we describe our spherical collapse model in Appendix~\ref{sec:SC}. Before we proceed, let us briefly discuss the connection with some of the literature on the subject. The halo mass function for a Yukawa theory like the one we study was computed by \cite{Martino:2009}. Our paper follows their formalism, and it is in a sense a straightforward extension to compute the conditional mass function and halo bias. Halo bias in $f(R)$ gravity and the DGP models has been measured from numerical simulations in \cite{Schmidt:2009,Khoury2009,Schmidt2009dgp,Chan2009,Schmidt:2010} . They find fair agreement between simulations and the bias derived from a modified Sheth and Tormen \cite{Sheth:1999mn} mass function whose parameters reflect the altered spherical collapse. There is a large literature on testing GR using the growth of large-scale structure, e.g. \cite{Lue2003}. Most studies allow for a scale-dependent growth factor, but ignore its effect on the galaxy bias. Our expression for the bias should be useful for incorporating the latter effect into such studies. There is also a substantial literature on a large-scale scale-dependent bias from primordial non-Gaussianity, e.g. \cite{Dalal2008}. This other source of scale dependence should be distinguishable from that from growth; we will discuss this in Sec. \ref{sec:conclusions}.

Conclusions} We have worked out the large-scale halo bias $b_1$ in a theory that has a scale-dependent linear growth, such as modified gravity or clustered dark energy. We give several expressions in \S \ref{sec:bias} \& \ref{sec:flatbarrier}. The most general expression is Eq. (\ref{eq:b}). It relates $b_1$ to the first crossing distribution $f$, which can be computed using the method described in \S \ref{sec:calcf}, for a general curved barrier. Note that all quantities describing the random walk: $\delta_1, S_1, \delta_0, S_0$ are defined at the initial redshift $z_i$, as opposed to the formation redshift $z_{\rm form}$. A much simpler expression, Eq. (\ref{formationbiastake2}), obtains in cases where the barrier can be approximated as flat. As can be seen in the example depicted in FIG. \ref{fig:barrier}, the barrier is generally not flat once a scale-dependent growth factor is allowed. However, it is often true that the barrier deviates from flatness only for the largest masses, and a flat barrier approximation (set at the small mass level) actually works quite well in predicting the halo bias for more modest halo masses (FIG. \ref{fig:bratio}). The expression in Eq. (\ref{formationbiastake2}) is particularly easy to use because the threshold for collapse $\delta_{1, {\rm form}}$ and the variance $S_{1,{\rm form}}$ are defined at the formation redshift as usual. Our main result is the imprinting of a scale dependence on the formation bias, by the modified linear growth factor. This can be understood to occur because regions of the same overdensity but different sizes today would not all have had the same overdensity at an earlier redshift, and hence would have varying halo densities (see Sec.~\ref{sec:birth}). The scale dependence of the growth factor manifests itself at a single redshift, making it possible to look for it with a local-universe snapshot. Using this effect and an appropriate galaxy sample, one can place constraints on gravity theories, as we demonstrate with a simple Yukawa theory. We have focused on the case in which the formation bias is deterministic. However, this bias may be stochastic, in which case this stochasticity will also be scale-dependent \cite{Hui:2007zh}. Although we will not perform a detailed analysis of stochastic bias here, we note that, even if the formation bias is deterministic in, say, $k$-space, the scale dependence means that it will be stochastic in real space (and vice-versa) \cite{Desjacques:2010}. Thus, the stochasticity of the bias is another potential signature of modified gravity theories. Looking forward, it would be interesting to work out how a scale-dependent large-scale bias can be incorporated into optimal weighting schemes for recovering the mass distribution from a galaxy or halo catalog \cite{Hamaus2010,Cai2010}.

1012.1335

1012

1012.4510_arXiv.txt

Soft gamma repeaters and anomalous X-ray pulsars are believed to be magnetars, i.e. neutron stars powered by extreme magnetic fields, $B\sim10^{14}$-$10^{15}$ Gauss. The recent discovery of a soft gamma repeater with low magnetic field ($< 7.5\times 10^{12}$\,Gauss), \src , which shows bursts similar to those of SGRs, implies that a high surface dipolar magnetic field might not be necessary for magnetar-like activity. We show that the quiescent and bursting properties of \src~ find natural explanations in the context of low-magnetic field Quark-Nova (detonative transition from a neutron star to a quark star) remnants, i.e. an old quark star surrounded by degenerate (iron-rich) Keplerian ring/debris ejected during the Quark-Nova explosion. We find that a 16 Myr old quark star surrounded by a $\sim 10^{-10}M_{\odot}$ ring, extending in radius from $\sim 30$ km to $60$ km, reproduces many observed properties of \src. The SGR-like burst is caused by magnetic penetration of the inner part of the ring and subsequent accretion. Radiation feedback results in months-long accretion from the ring's non-degenerate atmosphere which matches well the observed decay phase. We make specific predictions (such as an accretion glitch of $\Delta P/P \sim - 2\times 10^{-11}$ during burst and a sub-keV proton cyclotron line from the ring) that can be tested by sensitive observations.

Soft $\gamma$-ray Repeaters (SGRs) are sources of recurrent, short ($t \sim 0.1\,\mathrm{s}$), intense ($L \sim 10^{37-42}~\rm{ergs}$) bursts of $\gamma$-ray emission with an energy spectrum characterized by temperatures of $\sim 100$ keV. Occasionally SGRs enter into active episodes producing many short X-ray bursts; extremely rarely (about once per 50 years per source), SGRs emit a giant flare, an event with total energy at least 1000 times higher than their typical bursts. The normal pattern of SGRs is intense activity periods which can last weeks or months, separated by quiescent phases lasting years or decades. Current theory explains this energy release as the result of a catastrophic reconfiguration of a magnetar's magnetic field. AXPs are similar in nature but with a somewhat weaker intensity and no recurrent bursting. Several SGRs/AXPs have been found to be X-ray pulsars with unusually high spin-down rates, usually attributed to magnetic braking caused by their super-strong magnetic field. In all sources with magnetar-like activity, the dipolar fields span $5\times10^{13}\,{\rm G} <B< 2\times10^{15}\,{\rm G}$, which is $\sim 10$--1000 times the average value in radio pulsars. Magnetar-like activity previously was observed only in sources with dipolar magnetic fields stronger than the electron quantum field, $B_{\rm Q}=m_e^2c^3/e\hbar\sim 4.4\times10^{13}$ Gauss. \src\, was discovered on 5 June 2009 when the Fermi Gamma-ray Burst Monitor (GBM) observed two magnetar-like bursts (\cite{vanderhorst10}). Follow-up observations with several x-ray satellites show that it has x-ray pulsations at $\sim$9.1\,s, well within the range of periods of SGR sources (\cite{gogus09,esposito10}). The implied upper limit on the period derivative of \src\ of $\dot{P} < 6.0\times10^{-15}$\ss\ is by far the smallest of all known SGRs/AXPs. The corresponding limit on the surface dipolar magnetic field of \src\, is $B < 7.5\times10^{12}$\,Gauss, making it the SGR with the lowest surface dipolar magnetic field yet. The upper limit on the period derivative implies a characteristic age of the source in excess of $ > 24$\, Myr in the standard dipole model, $P/(2\dot{P})$. Despite such a low surface magnetic field ($< <B_{\rm Q}$), \src\ exhibits all the typical characteristics of an SGR. Its bursting properties can be summarized as follows (\cite{vanderhorst10,gogus09,esposito10}): \begin{itemize} \item At a distance of 2 kpc, assuming that the source is located in the Perseus arm, the estimate energies of the two observed bursts (in the 8-200 keV range) are $4\times 10^{37}$ erg and $ 2\times 10^{37}$ erg, which is at the lower end of the distribution compared to other SGR bursts but at the high end for AXP ones. \item An optically-thin thermal bremsstrahlung provides the best fit in both bursts (see Table 2 in \cite{vanderhorst10}). The spectrum softens from $kT\sim 33$ keV in the frist burst to $kT\sim 20$ keV in the second burst in the time period of $\sim 20$ minutes which separated the two bursts. \item Immediately following the burst and for the first 160 days (before it disappeared behind the sun), \src ~ flux declined by an order of magnitude from $\sim 3\times 10^{-11}$ erg cm$^{-2}$ s$^{-1}$ ($L_{\rm X}\sim1.4\times 10^{34}$ erg s$^{-1}$ ) to $\sim 3\times 10^{-12}$ erg cm$^{-2}$ s$^{-1}$ ($L_{\rm X}\sim 1.4\times 10^{33}$ erg s$^{-1}$). The corresponding blackbody temperature declined from 1 keV to 0.8 keV (see Figure 2 in \cite{esposito10}). \item The burst luminosity is $\sim 10^{39}$ erg s$^{-1}$ with burst temperature in the 20-30 keV range. Assuming blackbody emission the emitting area can be estimated to be on average $A_{\rm b}\sim 8.5\times 10^8$ cm$^2$. \end{itemize} \src's~ current properties (i.e. following the bursting era) show a spectrum which is well fit by an absorbed blackbody with a line-of-sight absorption \nh$=(1.5\pm1.0)\times10^{21}$\cm2 \, and $kT=0.67\pm0.11$\,keV. Using the current luminosity of $6\times 10^{31}$ erg s$^{-1}$, and a blackbody temperature of 0.67 keV (\cite{rea2010}) the emitting area is $A_{\rm q}\sim 3.1\times 10^{8}$ cm$^2$ ($<< 4\pi (10\ {\rm km})^2$) indicative of a hot spot on the surface of the star. Since the internal field strength required to produce crustal cracking in the Magnetar model should be typically in excess of $10^{14}$\, Gauss (\cite{td95}), one wonders how \src, with its weak surface magnetic field, can harbor a much stronger magnetic field in its crust. In fact, for such an old source ($> 24$ Myr) it would take an even stronger internal field to crack the cold crust\footnote{However there are no observational constraints on multipole moments of the surface field or on the internal toroidal magnetic field to rule out the crust cracking model. The tearing mode instability model proposed for magnetic field decay could account for short times scale bursts ($< 1$ yr) if the dissipation scale is much shorter than typical crust scale ($<< 10^4$ cm) (\cite{lyutikov2003}).}. Could such a difference between the surface and crustal field be possible and sustainable? What mechanism could lead to such a gradient, and if it exists does it mean that some ordinary pulsars are dormant SGRs waiting to erupt? Maybe a strong magnetic field is not necessary to explain SGR behavior. The existence of radio pulsars with $B > B_{\rm Q}$ and, so far, showing only normal behavior (\cite{kaspi10}) is another clue that magnetic fields larger than the quantum electron field alone may not be a sufficient condition for the onset of magnetar-like activity. Here we present an alternative model which offers natural answers to these outstanding questions. In our model high magnetic field strength is not necessary to explain the bursting phase of SGRs and AXPs : It involves an aligned rotator (a quark star; hereafter QS) and iron-rich debris material in the close vicinity ($\sim 20$-$100$ km) of the QS star. The QS is the compact remnant of the Quark-Nova (QN) explosion, a detonative transition from a neutron star (NS) to a QS (\cite{ouyed2002,vogt2004,niebergal2010b}). The QN detonation also leads to ejection of the NS outer layers (\cite{keranen2005}). If the QS is born slowly rotating, then the debris formed from the QN ejecta will be in co-rotation with the star's dipole field (\cite{oln1}, herefafter OLNI). Sources born with faster rotation will confine the debris into a Keplerian ring at 20-100 km away from the star (\cite{oln2}, hereafter OLNII)\footnote{This ring is unlike a fall-back disk around a neutron star (e.g. \cite{trumper2010}). The ring is iron-rich, very close to the star and degenerate. Similar ring formation when a neutron star is born appears implausible since the proto-neutron star is too large. Later on, there is no mechanism to eject degenerate material unless a violent change of state, like a QN occurs.}. In our model SGRs are QS with a co-rotating shell while AXPs are QS with a Keplerian ring. The debris consists of $\sim 10^{-6}M_{\odot}$ of iron-rich degenerate material. The initial QS surface magnetic field is $10^{15}$ G. Such initial extreme surface magnetic fields are natural values for quark stars experiencing color ferromagnetism before they enter the color superconducting phase (\cite{iwazaki2005}). The paper is organized as follows. Section 2 is a summary of our previous works and sets the stage for this work. Section 3 is the application of our model to evolved accreting QS-ring systems. The thermal feedback between the QS and the ring depends critically on the ring geometry, which evolves from thick to thin. For old sources, the resulting behavior is in a different regime than considered in previous papers. This leads to new forms of the equations, which we then fit to the observations. We conclude in Section 4.

\begin{enumerate}[1.] \item \src's current ring's mass is $\sim 2.5\times 10^{-10}M_{\odot}$ extending from $R_{\rm in}\sim 36$ km to $R_{\rm out}\sim 56$ km (see \ref{fig:illustration}). The ring's mass at birth is $m_{\rm ring}^0 > 9\times 10^{-7}M_{\odot}$ found from the $\tau_{\rm ring} > t_{\rm age}$ constraint. Given their degenerate nature the rings would have a very weak optical signature (unlike what would be expected from non-degenerate fall-back disks around neutron stars). \item \src's fit parameters give $m_{\rm w}\sim 3\times 10^{-16}M_{\odot}$. The wall is consumed on timescales given by eq.(22) in OLNII which yields a wall consumption time of $< 1$ hour for \src. \item Our model predicts a glitch, during the bursting phase (see eq.(29) in OLNII), of \begin{equation} \frac{\Delta P}{P}\sim -1.3\times 10^{-10} \frac{P_{10}R_{\rm in, 15}^{23/8}m_{\rm ring, -7}^{6/7}\mu_{\rm q, 3.3}^{15/8}}{I_{45}\eta_{0.1}^{5/4}t_{\rm Myr}^{1/2}}\ , \end{equation} where $I_{45}$ is the star's moment of inertia in units of $10^{45}$ gm cm$^2$. Using the best fit parameters for \src~ gives $\Delta P/P\sim -2\times 10^{-11}$. \item Fundamentally, the two-components model (QS and ring) provides a natural explanation for multiple emission components: vortex annihilation, the HS on the QS and the illuminated ring. However, as the system ages (i.e. at smaller $\dot{P}$), the ring becomes thinner in height and more extended in radial width and area. The feedback effect is reduced and so the ring's temperature decreases which makes the ring's contribution during quiescence harder to detect for old sources. In the case of \src~, and after full recovery from bursting phase, we predict a 2-component spectrum. The HS at $\sim 0.67$ keV and the ring at $T_{\rm ring}\sim 6.3$ eV. A third, much weaker, component would be related to the emission from the magnetic field reconnection at the surface of the QS following vortex expulsion and annihilation. Assuming a BB emission, the temperature of the QS would be $T_{\rm QS}\sim 17.3\ {\rm eV}~ \dot{P}_{-11}^{1/2}/R_{\rm QS, 10}^{1/2}$ or $< 0.4$ eV in the case of \src. \item During the bursting phase the ring's atmosphere material (in particular protons and $\alpha$ particles) will emit a cyclotron line $2\pi \nu_{\rm p}\sim e B_{\rm in}/(m_{\rm H} c)$ with $B_{\rm in}= B_{\rm s} (R_{\rm QS}/R_{\rm in})^3$ and $B_{\rm s}=\sqrt{3\kappa P \dot{P}} \sim 5.2\times 10^{14} P_{10}^{1/2} \dot{P}_{-11}^{1/2}$. Recalling that a 1 keV line corresponds to $2.42\times 10^{17}$ Hz this gives $\nu_{\rm p}\sim 1\ {\rm keV} P_{10}^{1/2} \dot{P}_{-11}^{1/2} R_{\rm in, 15}^{-1/2}$. For \src~ we get $\nu_{\rm p}\le 0.05$ keV and $\nu_{\alpha}= 0.5\nu_{\rm p}\le 0.025$ keV. \item During quiescence, the ring atmosphere (and thus the accreted material), is composed mostly of pure iron group nuclei. However, during the bursting phase and after irradiation, the atmosphere should be composed mostly of protons, of $\alpha$ particles and of ionized nuclei with $A \sim 28$ ($\mu=\mu_{\rm b}\sim 1.5$). It would be interesting to look for such signatures in emission during quiescence and bursting episodes. \item A key difference between objects showing magnetar-like activity and regular pulsars is the lack of persistent radio emission in the former class of objects. In our model the vortices force the interior magnetic field to align with the rotation axis, thus inhibiting persistent radio pulsation (\cite{ouyed2004,ouyed2006}). However, if \src~ has entered the last stages of ring consumption ($t_{\rm age}\sim \tau_{\rm ring}$), it should eventually show sporadic radio emission as it makes its way back to the vortex band as a RRAT (see OLNIII). This suggests that the system is in the last $\sim 1$\% of its lifetime and currently descending from the accretion band back to the vortex band. \end{enumerate} The above listed predictions in general produce weak signals: (i) Predictions 1\&4 yield a flux of $\sim 10^{-19}$ ergs cm$^{-2}$ s$^{-1}$ $\AA^{-1}$ which corresponds to a 26-27 magnitude in the V-band at 2kpc. This is in principle detectable from eight meter class telescopes. The signal from vortex annihilation is much weaker; (ii) for prediction 2, the next burst is not expected until of order of 2000 years from now; (iii) for 3 the glitch is small compared to ordinary pulsar glitches ($10^{-9}$ to $10^{-7}$) requiring timing precision higher than possible for a transient source; (iv) for 5 the frequency is below the Lyman limit and the line would be absorbed by the hydrogen in the ISM; (v) for 6, the flux during outburst is high enough to check existing X-ray spectra for intermediate elements lines; for 7, unfortunately the timescale to evolve into a RRAT is of the order of $10^4$-$10^5$ years. In summary, predictions 4 and 6 are the most testable.

1012.4510

1012

1012.1273_arXiv.txt

The Larkin-Migdal approach to a cold superfluid Fermi liquid is generalized for a non-equilibrium system. The Schwinger-Keldysh diagram technique is applied. The developed formalism is applicable to the pairing in the states with arbitrary angular momenta. We consider the white body radiation problem by calculating probabilities of different direct reactions from a piece of a fermion superfluid. The closed diagram technique is formulated in terms of the full Green's functions for systems with the pairing correlation. The cutting rules are used to classify the diagrams representing one-nucleon, two-nucleon, etc. processes in the matter. The important role of multi-piece diagrams for the vector-current conservation is demonstrated. In the case of equilibrated systems, dealing with dressed Green's functions, we demonstrate correspondence between calculations in the Schwinger-Kadanoff-Baym-Keldysh formalism and the ordinary Matsubara technique. As an example we consider neutrino radiation from the neutron pair breaking and formation processes in case of a singlet pairing. Necessary correlation effects are included. The in-medium renormalization of normal and anomalous vertices is performed.

\label{sec:intor} \subsection{Historical remarks} \label{subsec:history} The phenomenological theory of normal Fermi liquids at zero temperature was proposed by L.D. Landau in Refs.~\cite{Landau58,LP1981}. A.B. Migdal made the very important observation that a jump in the particle momentum distribution at the Fermi momentum corresponds to a pole of the fermion Green's function in the normal Fermi liquid and is preserved even in the strongly interacting system \cite{Mjump}. The presence of the pole contribution to the fermion Green's function allowed Galitsky and Migdal to develop the Green's function formalism for many-body fermionic systems, see Ref.~\cite{Mqp}. These concepts were first elaborated on example of low-lying particle-hole excitations in Fermi liquids. A.B. Migdal was first who applied these methods to description of various nuclear phenomena and constructed a closed semi-microscopic approach that is usually called "Theory of finite Fermi systems"~\cite{M67a,M67}. A general understanding of the phenomenon of superconductivity in case of the weak attraction between fermions was achieved by J.~Bardeen, L.N.~Cooper and J.R.~Schriffer in Ref.~\cite{BCS57}, see also Ref.~\cite{Schriffer} for detailed exposition. Due to the sharpness of the Fermi surface provided by presence of the Migdal's jump one can consider the fermions on the Fermi surface as moving in an effective two-dimensional momentum space. It follows immediately that even a weak attraction between two particles is sufficient to form a Cooper pair. As soon as the pairing phenomenon is established one can follow different routes in description of the superconductivity and superfluidity: In Ref.~\cite{Bog58} N.N.~Bogolyubov suggested a very convenient transformation from the particle $\psi$ operators to the new operators of effective excitations on top of the background of Cooper pairs. This transformation is broadly used in the theory of superconductivity. L.P.~Gorkov developed the Green's function formalism for superconducting fermion systems with an electron-phonon interaction~\cite{Gor58}. Y.~Nambu introduced a matrix formalism to the theory of superconducting metals \cite{Nam60} (Green's functions formulated in the Nambu-Gorkov space). In Ref.~\cite{El60} G.M.~Eliashberg extended the A.B.~Migdal's theory of the strong electron-phonon interaction in normal metals~\cite{MigEl} to include the Cooper pairs. This approach can be used to describe strong coupling superconductors. Also A.B.~Migdal was the first who rose the idea about a possibility of the neutron-neutron pairing and superfluidity in neutron stars, which where hypothetical objects that time~\cite{M59}. The Fermi liquid theory was then generalized by A.I.~Larkin and A.B.~Migdal for the description of fermion superfluids at zero temperature~\cite{LM63}. Their formulation is more general than that done in papers by Y. Nambu and L.P.~Gorkov, since it allows for different interactions in the particle-particle and particle-hole channels. A.J.~Leggett applied this formalism for the superfluid $^3$He at a finite temperature~\cite{Leg65a,Leg65}. J.~Schwinger in Ref.~\cite{Schw61}, L.P.~Kadanoff and G.~Baym in Ref.~\cite{KB62}, and L.P.~Keldysh in Ref.~\cite{Kel64} developed the non-equilibrium diagram technique for the description of non-equilibrium Fermi and Bose systems. Even for equilibrium systems at $T\neq 0$ Schwinger-Kadanoff-Baym-Keldysh approach is in many cases more convenient than the standard Matsubara technique (applicable only for equilibrium systems) since it does not involve the Wick rotation and the obtained results can be continuously transformed to those computed in the standard Feynman-diagram technique at zero temperature. The importance of coherence time effects on the production and absorption of field quanta from the motion of source particles in matter has been first discussed by L.D.~Landau and I.Ya.~Pomeranchuk~\cite{LandauP}. In Ref.~\cite{LPM} A.B.~Migdal developed the complete theoretical framework for the description of the bremsstrahlung radiation of ultra-relativistic electrons in the process of multiple rescatterings on Coulomb centers. Successful measurements of such a suppression of the bremsstrahlung radiation have been recently carried out at the Stanford Linear Accelerator Center~\cite{eSLAC}, see also the review in Ref.~\cite{SKlein}. Now this effect is named the Landau-Pomeranchuk-Migdal effect. In the framework of his theory of finite Fermi systems A.B.~Migdal developed the description of the soft pion degree of freedom in nuclear matter in application to atomic nuclei and neutron stars. In vacuum, pions are the lightest quanta of the strong interactions between baryons. In medium, pionic modes are softened even further due to the coupling to nucleon particle-hole modes and can be easily excited even at low excitation energies, similar to phonons in solids. As an intriguing consequence of the pion softening A.B. Migdal suggested a possibility of the pion condensation\footnote{Independently pion condensation was also suggested by D.J. Scalapino and R. Sawyer \cite{SS72}.} at the increase of the baryon density, see Refs.~\cite{M71,M78,MSTV90}. Latter on, in analogy to the pion condensation, the ideas of the kaon condensation \cite{BLRT94} and the charged $\rho$ meson condensation \cite{V97,KV05} in the interiors of neutron stars were explored. Softening of the pionic mode at finite temperature~\cite{VM78,VM82,D82} and at non-equilibrium~\cite{VS87,V93,VBRS95} may manifest in neutron stars~\cite{V01} and heavy ion collisions \cite{V93,RW}. A.B. Migdal rose question on a possibility of existence of superdense abnormal nuclei glued by the pion condensate \cite{M71,MMMS77}. Also a possibility of nuclei-stars was considered in Ref.~\cite{VSC77}. The similar ideas on a possibility of quark nuclei, quark stars and hybrid stars \cite{W84,AFO} are continued to be extensively explored nowadays, see \cite{Bombaci,BKV00}. \subsection{White body radiation} Below we consider the white-body radiation from a piece of a superfluid fermion matter. To be specific we focus on the neutrino radiation from a piece of superfluid nucleon matter. Standard Feynman technique of summation of squared matrix elements of reactions fails to calculate reaction rates in the medium, since in general case there are no asymptotic states for source particles in matter. Indeed, source particles continue to collide before and after radiation of a quantum. This gives rise to finite imaginary parts of the self-energy functions (particle widths). If one naively replaced the summation of all perturbative Feynman diagrams (with free Green's functions) by the summation of corresponding diagrams with dressed Green's functions, it would lead to a double counting due to multiple repetitions of some processes (for an extensive discussion of how one can treat this defect see \cite{VS86,VS87,KV95,KV99}). This calls for a formalism dealing with closed diagrams (integrated over all possible in-medium particle states) with full non-equilibrium Green's functions. Such a general formalism was developed in Ref. \cite{KV95}. It treats on equal footing one-fermion and multi-fermion processes as well as resonance reaction contributions of the boson origin, such as processes with participation of zero sounds and reactions on the boson condensates. Decomposition of diagrams is done in terms of the full $G^{-+}$ Green's functions (Wigner densities). Each diagram in the series with full Green's functions is free from the infrared divergencies. In such a way one generalizes Landau-Pomeranchuk-Migdal treatment of the multiple scattering on external centers to the treatment of the multiple scattering in matter. Both, the correct quasi-particle and quasi-classical limits are recovered. The formulation of the radiation problem in terms of closed diagrams calculated within the non-equilibrium Green's function in quasi-particle approximation was performed in Ref.~\cite{VS87} This approach was called the {\em "optical theorem formalism"}. In Refs.~\cite{VS86,VS87} it was demonstrated that the standard calculations of reaction rates via integration of squared reaction matrix elements and the results of the optical theorem formalism match exactly, provided conditions for the quasi-particle approximation for fermions are fulfilled. Formally the matching is done by cutting the closed diagrams. In general case considered in Ref.~\cite{KV95} going beyond the quasi-particle approximation, the series of closed diagrams is constructed with respect to the number of the $G^{-+}$ Green's functions. For low temperatures each $G^{-+}$ line brings extra $(T/\epsilon_{\rm F})^2$ factor in the production rate of the radiating quanta, $\epsilon_{\rm F}$ is the Fermi energy. In Ref.~\cite{KV95} the relations between reaction rates at finite widths and the quasi-particle rates were found. All real calculations of fermion superfluids were performed within quasi-particle approximation for fermions (when fermion width is much less than all other relevant energetic scales). Below we focus on the Larkin-Migdal approach to the cold fermion superfluids and formulate it in terms of the Schwinger-Kadanoff-Baym-Keldysh technique to describe fermion superfluids in equilibrium at $T\neq 0$ and out of equilibrium. \subsection{Neutrino cooling of neutron stars} Physics of neutron star cooling is based on a number of ingredients, among which the neutrino emissivity of the high-density hadronic matter in the star core is the important one. After the first tens of seconds (at most hours), the typical temperature of a neutron star decreases below the so-called neutrino-opacity temperature $T_{\rm opac}\sim (1 - 2)$~MeV. At these conditions neutrinos and anti-neutrinos can be radiated directly from the star interiors without subsequent rescattering, since their mean-free path is much longer than the star radius~\cite{VS86}. Hence, the star can be considered as a piece of a warm ``white'' body for neutrinos. Typical averaged lepton energy ($\sim$ several $T$) is much larger than the nucleon particle width $\Gamma_N \sim T^2/\varepsilon_{{\rm F}}$. Therefore, the nucleons can be treated within the quasi-particle approximation. This observation simplifies consideration essentially. One usually follows an intuitive way for the separation of the processes according to their phase spaces. The one-nucleon processes (if they are not forbidden by the energy-momentum conservation) have the largest emissivity, $\epsilon_{\nu}\propto T^6$ for non-superfluid systems, then two-nucleon processes come into play, $\epsilon_{\nu}\propto T^8$, and so on. In the optical theorem formalism one-nucleon processes are determined by the self-energy $\Sigma^{-+}$ of virtual $W$ and $Z$ bosons expanded in the series with respect to the number, $N$, of $G^{-+}G^{+-}$ loops with full $"++"$ and $"--"$ vertices. The $N=1$ diagrams correspond to one-nucleon processes, the $N=2$ diagrams to the two-nucleon processes, etc. In the so-called {\em "standard scenario"} of the neutron star cooling, the processes were calculated without taking into account in-medium effects. It was argued that the most important channel at temperatures up to $T\sim 10^{8}$--$10^{9}$~K is the modified Urca (MU) process $n \, n \rightarrow n \, p\, e\, \bar \nu$. First estimates of the MU emissivity were done in \cite{BW65,TC65}. In Ref.~\cite{FM79,M79} B.~Friman and O.V.~Maxwell recalculated the emissivity of this process in the model, where the nucleon-nucleon (NN) interaction was approximated by a free one-pion exchange (FOPE). The expression of the neutrino emissivity obtained by them was used in various computer simulations, e.g., in Refs.~\cite{T79,NT81,SWWG96}. Besides the MU process, the {\em "standard scenario"} includes also the processes of the nucleon (neutron and proton) bremsstrahlung (NB) $n\, n\rightarrow n\, n \nu \bar{\nu}$ and $n\, p\rightarrow n\, p \nu \bar{\nu}$, which contributions to the emissivity is smaller than those of the MU processes, see Refs.~\cite{FSB75,FM79}. The density dependence of the reaction rates calculated with the FOPE is rather weak and the neutrino radiation from a neutron star depends very weakly on the star mass. There exists another class of so-called ``exotic'' processes, which occur only if some special condition is fulfilled, i.e. when the nucleon density exceeds some critical values. These are the direct Urca (DU) processes on nucleons (e.g., $n\rightarrow pe\bar{\nu}$) and hyperons \cite{LPPH91}, pion Urca reactions on a pion condensate \cite{MBCDM77,VS84,VS86}, kaon Urca processes on a kaon condensate \cite{BKPP88,T88}, $\rho$-Urca processes on charged $\rho$-condensates \cite{V97,KV05}, DU processes on quarks \cite{Iwamoto}, DU processes on fermion condensates \cite{VKZC00}. The values of critical densities are different for various processes and are model dependent. For example, some relativistic mean-field models produce the critical density of the DU reaction, $n_{c}^{\rm DU}$ as low as the nuclear saturation density $n_0 \simeq 0.16$~fm$^{-3}$. However, the realistic, microscopically-based Urbana-Argone equation of state~\cite{APR98} yields $n_{c}^{\rm DU}\simeq 5n_0$. The simulations of the neutron star cooling history in Refs.~\cite{BGV,GV} have shown that the occurrence of the DU processes in the neutron star with masses $M<1.5 M_{\odot}$ would lead to problems with the explanation of soft $X$-ray data. The constraint on the equation of state of the dense nuclear matter, requiring a sufficiently high value of $n_{c}^{\rm DU}$, was proposed in Ref.~\cite{KV05} and explored in details in Ref.~\cite{Army}. It was shown in Refs.~\cite{VS84,VS86,VS87,SV87,MSTV90,V01} that the neutrino emission from dense hadronic component in neutron stars is subject of strong modifications due to collective effects in the nuclear matter. Many new reaction channels open up in medium in comparison to the vacuum. In Refs.~\cite{VS84,VS86,VS87,SV87,MSTV90,V01} the nucleon-nucleon interaction was considered within the Landau--Migdal approach to Fermi liquids. The softening of the in-medium one-pion exchange (MOPE) mode and other medium polarization effects, like nucleon-nucleon correlations in the vertices, renormalization of the local part of $NN$ interaction due to loop effects, as well as a possibility of the neutrino emission from intermediate reaction states and DU-like reactions involving zero sounds and boson condensates were incorporated. It was demonstrated in Refs.~\cite{VS86,SV87,MSTV90,V01}, that for $n \gsim n_0$ the neutrino emissivity is mainly determined by the medium modified Urca (MMU) process, in which the neutrino is radiated from the intermediate reaction states. This fact changes significantly the absolute value and the density dependence of the $nn\rightarrow npe\bar{\nu}$ process rate. The latter becomes very strong. Therefore, for neutron stars with larger masses the resulting emissivity of the MMU process proves to be substantially higher than the corresponding value (MU) calculated in the FOPE model of Ref.~\cite{FM79}. For $n\gsim n_0$, the medium-modified nucleon bremsstrahlung (MNB) processes yield a smaller contribution than MMU ones since the former does not include the neutrino radiation from intermediate states, e.g. from intermediate pion. However, the MNB processes are more efficient than the NB ones for such densities. Oppositely, for $n<n_0$ the in-medium effects can moderately suppress the two-nucleon reaction rates compared with those given by the FOPE model~\cite{BRSSV,HPR,Schwenk04}. The pion softening effect disappears at $n\lsim 0.5$--$0.7\,n_0$, see Refs.~\cite{MSTV90,V01}, but the nucleon-nucleon short-range repulsion effect remains. Inclusion of the nucleon-nulceon correlations without the pion softening~\cite{RPLP} yields a suppression effect also at $n\gsim n_0$. Obviously, this effect also follows from general consideration in Refs.~\cite{VS86,MSTV90}, if one artificially suppresses the pion softening effect. After the seminal work of A.B.~Migdal~\cite{M59}, various aspects of the nucleon superfluidity in neutron stars were studied in the literature: The presence of a nucleon superfluid interacting with the normal component is needed for explanation of glitches in pulsar periods and neutron star quakes~\cite{ST83}. Explanation of pulsar cooling curves also requires an inclusion of superfluid phases~\cite{GV}. Several superfluid phases are found possible. Phase transitions between different phases may take place~\cite{KCZ}. It is commonly accepted that most important are the superfluid phases with the spin-singlet pairing of neutrons and protons, in the $1S_0$ state, and the spin-triplet pairing of neutrons in the $3P_2$ states. The latter is believed to occur in neutron star interior at $n\gsim n_0$ in the state with $m_J =0$, where $m_J$ is the projection of the total pair momentum onto a quantization axis. In case $|m_J| =2$ the exponential suppression of the specific heat and the neutrino emissivity is replaced by a power-law suppression since the gap vanishes at the poles of the Fermi sphere. This possibility was mentioned for the first time in Ref.~\cite{VS87}, the corresponding reaction rates were calculated in Ref.~\cite{YKL99}. However, a mechanism to realize this interesting possibility in neutron star cooling was not elaborated yet. Many papers are devoted to the calculation of pairing gaps within different approaches~\cite{APW,Tamagaki70,Amundsen85,Takatsuka93,KKC96,Schulze96,Elgaroy98,Khodel01,SF,KCTZ,Hebeler07,Chen08}. The obtained results can be essentially different depending on a model for the nucleon-nucleon interaction and a calculation scheme. The predictions of the neutron $3P_2$ gaps are especially uncertain, e.g., compare Refs.~\cite{SF} and ~\cite{KCTZ}. For review see Ref.~\cite{SC06} and references there in. Ref.~\cite{SF} argues that $3P_2$ gap should be strongly suppressed whereas Ref.~\cite{KCTZ} argues for its strong enhancement. Reference~\cite{GV} calculated cooling curves using both these assumptions and concluded that the cooling history is naturally explained within assumption on the suppressed $3P_2$ gap. Recently Ref.~\cite{PPLS10} studied the new data on the cooling of Cas A object. Their conclusion is in favor of a suppressed $3P_2$ gap. At temperatures below the critical temperatures of the neutron, $T_{cn}$, and proton, $T_{cp}$, pairing, the reaction rates, considered above, are suppressed because of a decrease of the available phase space. Initially, the suppression effects were included simply by multiplying the rate of a two-nucleon process by the factors $e^{-2\Delta/T}$~\cite{M79}. Later, the phase-space suppression factors (so called $R$ factors) have been treated more accurately in Ref.~\cite{YLS99}. In nucleon superfluids, there exist new neutrino-production mechanisms, which are forbidden for $T>T_c$. These are the processes, suggested in Ref.~\cite{FRS76,VS87,SV87}, in which the creation of a neutrino--anti-neutrino pair is associated with the breaking and formation of a Cooper pair -- the so-called nucleon pair breaking and formation (PBF) processes. The emissivities of the nucleon PBF processes are suppressed at $T<T_c$ by the same factor $\sim \exp (-2\Delta/T)$ as for the MU, NB, MMU and MNB processes. However, in comparison to the all latter processes, the nucleon PBF processes have the large one-nucleon phase-space volume~\cite{VS87,SV87}. The existence of this new cooling mechanism demonstrates that influence of the nucleon pairing on the neutrino production rates cannot be reduced just to an introduction of a simple phase-space suppression factor. Early works~\cite{FRS76,VS87,SV87,V01,KHY,YKL99,YLS99} which studied the PBF processes, did not care about the conservation of the weak vector current. The latter is fulfilled only if the in-medium renormalization of weak vertex functions is performed in accord with the renormalization of Green's functions. This problem was tackled in Refs.~\cite{LP,KV1,KV2,KR,LPpairing,SMS,SR09}. Reference~\cite{LP} argued that the emissivity of the $1S_0$ PBF processes should be dramatically suppressed as $\propto v_{\rm F}^4$, where $v_{\rm F}$ is the Fermi velocity of non-relativistic nucleons, provided the vector current conservation constraint is taken into account. The consistent calculation of the PBF emissivity induced by the vector and axial-vector currents was performed in Ref.~\cite{KV1,KV2} within the Larkin-Migdal-Leggett Fermi-liquid approach. The latter takes properly into account correlation effects in both particle-particle and particle-hole channels. It was demonstrated that the neutrino emissivity is actually controlled by the axial-vector current and is suppressed only by the factor $\propto v_{\rm F}^2$, rather than $\propto v_{\rm F}^4$. Both neutron PBF and proton PBF processes yield contributions of the same order of magnitude provided strong and electromagnetic renormalizations of the proton weak vertices \cite{VS87,VKK,L00} are included. In Ref.~\cite{LPpairing} one argues that for the $3P_2$ neutron pairing the vector current conservation changes moderately the result obtained without its inclusion. As pointed out in Ref.~\cite{GBSMK}, the suppression of the PBF processes at low densities might be served as a possible explanation of the superburst ignition. As we have mentioned, the convenient Nambu-Gorkov formalism developed for the description of metallic superconductors, cf. Refs.~\cite{Nam60,Gor58,Schriffer}, does not distinguish interactions in particle-particle and particle-hole channels. These interactions can be, however, essentially different in a strongly interacting system, like in the nuclear matter and in the liquid He$^3$. The adequate methods for Fermi liquids with pairing were developed for zero temperature by A.I.~Larkin and A.B~Migdal in Ref.~\cite{LM63} (see also \cite{M67}) and for a finite temperature by A.J.~Leggett in Ref.~\cite{Leg65a,Leg65}. The problem of calculation of a response function of a Fermi system to an external interaction becomes tractable at cost of introduction of a set of Landau-Migdal parameters for quasi-particle interactions. Parameters can be either evaluated microscopically or extracted from analysis of experimental data, see Ref.~\cite{M67}. The technical difference of the Larkin-Migdal and Leggett approaches is that the former approach works out equations for full in-medium vertices, whereas the latter one calculates directly a response function. The former approach was aimed at the study of transitions in nuclei, and the latter on the analyzes of collective modes in superfluid Fermi liquid. The principal equivalence of both approaches was emphasized already by A.J.~Leggett in Refs.~\cite{Leg65a,Leg65}. Reference \cite{KV2} demonstrates how one may use both approaches in calculations of the PBF rates. Reference~\cite{SVSWW97} was the first one, in which the most important in-medium effects were incorporated in the numerical code for neutron star cooling. Among them neutron PBF and proton PBF processes were treated as equally important. The PBF processes (but with free vertices) were incorporated also in the ``standard'' cooling scenario \cite{P98,YLS99} that led the authors of Ref.~\cite{Page} to the suggestion of the minimal cooling paradigm. Detailed simulations of different medium effects have been done in \cite{BGV,GV}. In contrast to the minimal cooling paradigm, the medium modifications of all reaction rates lead to their pronounced density dependence. For the PBF processes it is mainly due to the dependence of pairing gaps and nucleon-nucleon correlation factors on the density. For MMU processes the reaction matrix elements are strongly density dependent due to the softening of the exchanged pion and the dependence of nucleon-nucleon correlation factors on the density. It establishes the strong link between the cooling behavior of a neutron star and its mass \cite{VS84,VS86,MSTV90,SVSWW97,BGV, GV}. The density dependence of the reaction rates provides a smooth transition from {\emph ``standard'' } to {\emph ``non-standard''} cooling for the increasing star-center density, i.e., for increasing the star mass. Thus, the inclusion of the most important in-medium effects within the {\em{ ``nuclear medium cooling scenario''}} enables us to describe appropriately both high and low surface temperatures obtained from analyzes of soft X-ray pulsar data. The mentioned above moderate suppression of the PBF emissivity ($\propto v_{\rm F}^2$) at $1S_0$ pairing should not significantly affect general conclusions on the neutron star cooling history done in previous works where it was not incorporated. The paper is organized as follows. In Section~\ref{sec:noneq} we formulate description of normal Fermi liquids at non-equilibrium. Softening of pionic degrees of freedom is taken into account. In Section~\ref{sec:LM} we perform generalizations to the fermion superfluids at non-equilibrium. The Larkin-Migdal equations are formulated on the Schwinger-Keldysh contour. A possibility of the pairing in an arbitrary momentum state is considered. In Section~\ref{sec:White} we introduce optical theorem formalism for normal and superfluid fermion systems out of equilibrium. Cutting rules for closed diagrams expanded in series with respect to the number of $G^{-+}$ full Green's functions are formulated. Important role of multi-piece diagrams is shown. Fermi liquid renormalizations are performed in Section~\ref{sec:Fermi-liquid}. Equilibrium $T\neq 0$ systems with pairing are considered in Section~\ref{sec:eqpairing}. In Section~\ref{sec:neutrino}, as an example, we find the current-current correlator and in Section \ref{sec:PBF}, the neutrino emissivity from the PBF processes on neutrons paired in $1S_0$ state. Technical details are given in Appendices.

A.I.~Larkin and A.B.~Migdal extended the Landau's Fermi-liquid theory onto superfluid systems. In this paper we re-formulated their approach for systems out of equilibrium. For that we used Schwinger-Kadanoff-Baym-Keldysh formalism. Important improvements of the Larkin-Migdal approach compared to the Nambu-Gorkov one are that the former approach allows to deal with strong interactions different in the particle-hole and particle-particle channels. These achievements have been used by A.J.~Leggett who generalized the Larkin-Migdal approach to describe strongly interacting fermion superfluids at finite temperatures and applied it to description of superfluid $^3$He. He used Matsubara diagram technique. The use of the Schwinger-Keldysh diagram technique allows to consider variety of non-equilibrium problems. In application to nucleon systems, in general, the considered in this paper formalism can be applied to the paring in the states with an arbitrary angular momentum; it operates with various forms of a nucleon-nucleon interaction: scalar, spin-spin, spin-orbit and tensor interactions. As argued by A.B.~Migdal the tensor forces mediated by the pion exchange should enhance (pion softening) with increase of the nucleon density. Inclusion of this effect might be important in the case of the $P$-pairing in neutron star interiors. We considered the neutrino radiation from a finite piece of the nuclear matter bearing in mind the problem of the neutron-star cooling. We used optical theorem formalism formulated in terms of closed diagrams with the full fermion and boson Green's functions and the full nucleon-nucleon interaction. The series of the diagrams is constructed with respect to the number, $N$, of the full $G^{-+}$ fermion Green's functions. For simplification we considered a system which evolves slowly in time and has small spatial gradients. This allowed us to perform the gradient expansion after the Wigner transformation and keep only gradient-independent terms in calculations of reaction rates. We demonstrated that in order to exactly satisfy the vector current conservation in the nucleon pair breaking and formation processes it is not sufficient to include only one $N=1$ term of the series, rather one needs to re-sum the RPA series including multi-piece diagrams. (The multi-piece diagrams decay in more than two pieces, being cut through $(-+)$, $(+-)$ lines). This demonstration shows, how one should separate one-nucleon, two-nucleon, etc. processes, in accordance with exact conservation law of the vector current. Comparison of the RPA $\Sigma^{-+}$ self-energy with the $N=1$ contribution shows the accuracy with which one may deal, using only one $N=1$ diagram. Then we demonstrated how the developed formalism allows to calculate neutrino emissivity from the piece of a warm nucleon matter in presence of the nucleon pairing. As simplest example we calculated neutrino emissivity in the neutron pair breaking and formation processes. These processes are of one-nucleon origin. To simplify consideration, we focused on the case of the ordinary $1S_0$ pairing of neutrons. More difficult is to calculate the emissivity of the two-nucleon ($N=2$) processes, and $N\geq 3$ processes in the presence of pairing. The existing nowadays results for the reaction rates in nucleon systems with pairing are based on the so-called $R$ phase-space suppression factors used to reduce the production rates calculated without pairing. Such an approach can be used only for rough estimations. The formalism formulated in the present paper is fully suited to properly perform the calculations. It is also interesting to search for new processes which might be open in the non-equilibrium and equilibrium medium because of the interaction between different reaction channels. These questions require a separate consideration. In the present paper we focused on the neutrino radiation problem. However, the white body radiation of other quanta can be considered in similar way. The calculated rate $\Sigma^{-+}$ can be considered as the gain term in the generalized kinetic equation for the virtual $W/Z$ boson or for the anti-neutrino. For consistency then one needs to include first-order gradient memory terms into the collision term. Another important question is how to go beyond the quasi-particle approximation for fermions in strongly interacting fermion systems with pairing. Within the quasi-particle approximation for fermions, the formalism based on the Fermi-liqiud renormalization is developed. To quantify the results it remains to know the Landau-Migdal parameters. For the problems under consideration one needs to know them as functions of the density, isospin composition, frequency and momentum. The information extracted from analysis of atomic nucleus experiments is definitely insufficient for these purposes. Existing calculations of the Landau-Migdal parameters are still incomplete. We hope that the present study will motivate further attempts to extract these parameters. In spite of all difficulties, it seems to be the cheapest way to achieve understanding of many new interesting phenomena occurring in the strongly interacting fermion systems in the presence of the pairing. The direction of the research was shown in the works done in 50th--70th years of the XXth century. The pioneering contribution to the development of the methods of the quantum many-body theory including the problem of fermion pairing belongs to Arkady~Migdal.

1012.1273

1012

1012.3089_arXiv.txt

We present the first white dwarf mass distributions of a large and homogeneous sample of post-common envelope binaries (PCEBs) and wide white dwarf-main sequence binaries (WDMS) directly obtained from observations. Both distributions are statistically independent, with PCEBs showing a clear concentration of systems towards the low-mass end of the distribution, and the white dwarf mass distribution of wide WDMS binaries being similar to those of single white dwarfs. Our results provide evidence that the majority of low-mass ($\Mwd \la 0.5\,\Msun$) white dwarfs are formed in close binaries.

The mass distribution of hydrogen atmosphere white dwarfs is strongly clustered around an average value of $\sim0.6\,\Msun$ \citep{koesteretal79-1, marshetal97-1, kepleretal07-1, holbergetal08-1}, as predicted by models of single star evolution. In addition to the pronounced peak at $\sim0.6\,\Msun$, a second peak at lower masses, $\sim0.4\,\Msun$, has been frequently found \citep[e.g.][]{bergeronetal92-1, bragagliaetal95-1, liebertetal05-1, kepleretal07-1}. Given the evolutionary time scale of low-mass single main sequence stars that are supposed to form such low-mass ($\Mwd \la 0.5\,\Msun$) white dwarfs significantly exceeds the Hubble time, the existence of low-mass white dwarfs has been interpreted as the result of strong mass transfer interactions in binaries. In this scenario the more massive (primary) star in a main sequence binary fills its Roche-lobe on the giant branch, and dynamical unstable mass transfer onto the less massive (secondary) star leads to the formation of a common envelope engulfing both the core of the primary star and the lower mass secondary star. Orbital energy released due to the shrinkage of the binary orbit is supposed to rapidly expel the envelope, and may terminate the growth of the He core of the primary star before it reaches a sufficient mass for He-ignition. The outcome of this close binary evolution is a post-common envelope binary (PCEB) consistent of a (possibly low-mass He-core) white dwarf and the (basically unaltered) low-mass secondary \citep{paczynski76-1, webbink84-1, iben+tutukov86-1}. The hypothetical binary origin of low-mass white dwarfs is observationally supported by the large fraction of low-mass white dwarfs in short orbital period double white dwarfs \citep[e.g.][]{marshetal95-1}, in neutron star-white dwarf binaries \citep[e.g.][]{sigurdssonetal03-1} and in subdwarf B star-white dwarf pairs \citep[e.g.][]{maxtedetal02-1}, as well as the existence of some PCEBs with low-mass white dwarf primaries \citep[e.g.][]{schreiber+gaensicke03-1, nebotetal09-1, pyrzasetal09-1, zorotovicetal10-1}. However, also a fair number of apparently single low-mass white dwarfs are known that exhibit neither radial velocity variations, nor infrared flux excess, the typical hallmarks of white dwarfs with close companions \citep{maxtedetal00-4, napiwotzkietal07-1, kilicetal10-1}. Possible explanations for the existence of these systems include supernova type Ia explosions in semi-detached close binaries that blow away the envelope of the companion thus exposing its low-mass core \citep{justhametal09-1}, severe mass-loss on the first giant branch \citep{kilicetal07-1} that may even lead to the formation of low-mass white dwarfs with CO-cores \citep{prada-moroni+straniero09-1}, stellar envelope ejection due to the spiral-in of close giant planets \citep{nelemans+tauris98-1}, or the merging of two very low-mass white dwarfs \citep{hanetal02-1}. So far, conclusive studies testing the close binary origin of low-mass white dwarfs have been prevented by the lack of a sufficiently large and homogeneous sample of PCEBs. This is now rapidly changing thanks to the Sloan Digital Sky Survey \citep[SDSS,][]{yorketal00-1, abazajianetal09-1, yannyetal09-1} from which $\sim2000$ white dwarf/main sequence (WDMS) binaries have been spectroscopically identified \citep{silvestrietal07-1, schreiberetal07-1, helleretal09-1, rebassa-mansergasetal10-1}. This population of WDMS binaries consists of wide systems whose stellar components did not interact and thus evolved like single stars, and close binaries that suffered from dynamically unstable mass transfer, i.e. PCEBs. Based on extensive radial velocity follow-up of the SDSS WDMS binary sample \citep{rebassa-mansergasetal07-1, rebassa-mansergasetal08-1, schreiberetal08-1, schreiberetal10-1}, we demonstrate here that the mass distribution of the white dwarfs in PCEBs does indeed differ significantly from that of white dwarfs that do not undergo binary interactions, containing a large fraction of low-mass white dwarfs.

\label{s-concl} The white dwarf mass distributions of PCEB and wide WDMS binary candidates are significantly different. While the mass distribution of wide WDMS binary candidates resembles those of single white dwarfs, the PCEB white dwarf mass distribution contains a large number of low-mass white dwarfs. Taking into account both the PCEB detection probabilities of the measurements and selection effects in SDSS we find that the large majority of low-mass white dwarfs resides in close binary stars. This result confirms a crucial prediction of current theories of close binary evolution and provides the so far strongest observational evidence for common envelope evolution forming most close compact binary stars. In agreement with the fraction of $\sim4$ per cent apparently single low-mass white dwarfs identified in single white dwarf samples, $\sim0.5-10$ per cent of the wide binaries in our sample seem to contain low-mass white dwarfs. These low-mass white dwarfs in wide binaries must have either formed due to exceptionally strong mass-loss of the primary progenitor on the first giant branch or are the descendants of triple systems, and some are expected to contain a close double degenerate primary component.

1012.3089

1012

1012.1874_arXiv.txt

We review computational approaches to understanding the origin of the Initial Mass Function (IMF) during the formation of star clusters. We examine the role of turbulence, gravity and accretion, equations of state, and magnetic fields in producing the distribution of core masses - the Core Mass Function (CMF). Observations show that the CMF is similar in form to the IMF. We focus on feedback processes such as stellar dynamics, radiation, and outflows can reduce the accreted mass to give rise to the IMF. Numerical work suggests that filamentary accretion may play a key role in the origin of the IMF.

The Initial Mass Function (IMF) plays a central role in astrophysics because it encapsulates the complex physics of star formation. Observations suggest that the IMF can be variously described as a piece-wise power-law \citep{Kroupa:2002}, a lognormal \citep{Miller/Scalo:1979}, or a lognormal distribution with a high-mass power-law tail \citep{Chabrier:2003}. The high mass behaviour of the IMF \cite{Salpeter:1953} is a power-law, $dN \propto m^{-2.3} dm$ for stellar masses $m \ge 0.5 M_{odot}$. The peak mass for isolated stars in the galactic disk is $ 0.1 M_{\odot}$ and $ 0.2-0.3 M_{\odot}$ \citep{Chabrier:2003} for the bulge. The form of the IMF is similar in many different galactic and extragalactic environments such as globular clusters, wherein one has a large range of metallicities and concentrations \citep{Paresce/DeMarchi:2000}. The current evidence therefore tends to support the notion that the IMF is universal. What physical processes produce the IMF? Observations show that it emerges during the early stages of the formation of star clusters \citep{Meyer/etal:2000,Zinnecker/etal:1993}. Young stars are formed within gravitationally bound subunits of a cluster-forming environment known as "cores" whose mass distribution - the core mass function (CMF) - strongly resembles the IMF. Cores are closely associated with filaments, as recent observations using the Spitzer and Herschel observatories clearly show \citep{Andre/etal:2010}. Competing physical processes such as turbulence, gravity, cooling and thermodynamics, as well as magnetic fields play significant roles in building the CMF, as well as filaments, within molecular clouds. Feedback processes such as radiation from massive stars, jets and outflows, as well as stellar dynamics serve to truncate the accretion of material onto stars and their natal disks. These may lead to the emergence of the IMF from the CMF. This review focuses on the critical role that computation is playing in exploring the origin of the IMF in clusters. We focus first on the processes leading to the CMF, and then discuss feedback processes that may convert the CMF into the IMF. Recent reviews of the computational aspects of the IMF may be found in \citet{MacLow/Klessen:2004}, \citet{Bonnell/etal:2007}, \citet{Larson:2007}, and \citet{Klessen/etal:2009}. The theory of the IMF is covered by Hennebelle (this volume), and \citet{McKee/Ostriker:2007}.

1012.1874

1012

1012.5982_arXiv.txt

A simple model which exhibits dynamical flame properties in 1D is presented. It is investigated analytically and numerically. The results are applicable to problems of flame propagation in supernovae Ia.

\label{sec:model} Our model includes two dependent variables: the temperature $T$ and the deficient reagent fraction $c$. There are also two independent variables: time $t$ and distance $x$. Neglecting diffusion of the deficient reagent (which is reasonable, because the kinetic coefficients in a white dwarf are such that $\mbox{\sl Le}\gg1$), we introduce the following system of equations for $T$ and $c$: \begin{equation} C\partial_t T=\kappa \partial_x^2 T+W\omega c\Theta(T-T_0),~~\partial_t c=-\omega c\Theta(T-T_0), \label{Ph010} \end{equation} where $\Theta(\dots)$ is the step-function. These equations describe deflagration burning in solid propellants, because two main physical processes which drive a slow front are present here: the thermal conductivity and the burning itself. The equations include the following constants: $C$ is the thermal capacity of the fuel per unit volume; $W$ is energy per unit volume of the deficient reagent fraction, released due to the reaction; and $\omega$ is the reaction rate, with dimensions 1/s. Assuming that the deficient reagent fraction before ignition is equal to $c_0$, one can see that the fuel temperature after complete burnout will increase by the value \begin{equation} T_f=\frac{Wc_0}{C}\, . \label{Ph020} \end{equation} Assuming further that the temperature $T(0)$ before ignition is much lower than $T_f$ we will neglect the former one, setting $T(0)=0$. In real white dwarfs, the thermal capacity depends on the temperature: for high density $\rho\sim 10^8\div 10^9$ g/cm$^3$ all parameters are determined by the relativistic degenerate electrons, that is $C\propto T$. But here we omit the dependence for the sake of possibility of analytical analysis. The latter concerns also the coefficient of thermal conduction $\kappa$ which is determined by both electrons and photons. The step wise approximation of the temperature dependence of nuclear reactions rates are used to have linear equations in both domains $T<T_0$ and $T>T_0$. Intrinsic nonlinearity of the problem is moved only to boundary conditions between the two regions. This simplification will allow us to investigate the problem analytically. However we believe that our approximation of the temperature dependence represents to some extend very sharp temperature dependence of nuclear reaction rates, especially for high values of the Zeldovich number: {\sl Ze}$\gg 1$. See Fig.~\ref{Ze-fig}. \begin{figure} \includegraphics[width=3 in, clip=]{Ze2.eps} \caption{Comparison of Arrhenius law for reaction rate and our step wise function for the same {\sl Ze}~$=9$.} \label{Ze-fig} \end{figure} Usually the Zeldovich number {\sl Ze} is defined assuming Arrhenius law for the temperature dependence of reactions rates: $w\propto \exp (-T_a/T)$. In this case the Zeldovich number is defined usually as: {\sl Ze}$=T_a/T_f$. To approximate this temperature dependence we may equate the temperatures for both dependence where the rates become $e$-fold lower than their maximum values at $T=T_f$. Then we will have the following definition of $T_0$ through $T_a$ and $T_f$: \begin{equation} {\sl Ze}=\frac{T_a}{T_f}=\frac{T_0}{T_f-T_0};\quad\mbox{~~or~~} T_0=T_f\, \frac{T_a}{T_f+T_a}=T_f\, \frac{{\sl Ze}}{1+{\sl Ze}} \, . \label{Ze} \end{equation} Applying the same procedure to the power law for the reaction rate ($w\propto T^n$), we would obtain the following relationship: $T_0=T_fe^{-1/n}\approx T_f\, (1-n^{-1})$ (for $n\gg 1$). We may introduce now new dimensionless variables labelled by ``$\tilde{\phantom{a}}$'' and defined as follows: \begin{equation} t=\tau \tilde{t}\, ;\quad x=l\tilde{x}\, ;\quad c=c_0\tilde{c}\, ;\quad T=T_f\tilde{T}\, , \label{Ph030} \end{equation} where \begin{equation} l=\sqrt{\frac{\kappa}{C\omega}\,\left(\frac{T_f}{T_0}-1\right)}\, ;\quad \tau=\frac{Cl^2}{\kappa}\, . \label{Ph040} \end{equation} Scales for $x$, $t$, $T$, $c$ are chosen for all characteristic parameters of the traveling wave in dimensionless units to be equal to 1. This concerns the velocity of the traveling wave, the concentration $c$ far before the wave, the temperature $T$ far beyond the wave, and the characteristic width of the wave. We will use below only the dimensionless variables skipping the tilde ``$\tilde{\phantom{a}}$''. Thus our dimensionless system is as follows: \begin{equation} \partial_t T=\partial_x^2 T+\omega_0 c\Theta(T-T_0),~~\partial_t c=-\omega_0 c\Theta(T-T_0)\, , \label{sys:simpl} \end{equation} where \begin{equation} \omega_0=\omega\tau = T_0/(1-T_0); \quad (0<T_0<1)\, . \label{T0} \end{equation} Dimensionless value of the igniting temperature $T_0$ is used in Eq.~(\ref{T0}). In initial physical units it is equal to $T_0/T_f$. Our effective Zeldovich number [Eq.~(\ref{Ze})] in dimensionless units can be expressed as follows: \begin{equation} \mbox{\sl Ze}=\omega_0=\frac{T_0}{1-T_0}\, . \label{Ze_m} \end{equation} The stationary traveling wave depends on $x$ and $t$ only via the combination $\xi=x-vt$. For this reason it is convenient to introduce new spacial coordinates $\xi=x-vt$ instead of $x$, where $v$ is the constant velocity of moving frame. Thus $T$ and $c$ depend now on $t$ and $\xi$. It is important that the system~(\ref{sys:simpl}) is linear in both $T>T_0$ and $T<T_0$. Nonlinearity appears only at the points, where the transition $0\rightarrow 1$ in $\Theta$-function occur. Our equations~(\ref{sys:simpl}) mean that $T$, $\partial_x T$ and $c$ are continues at $T=T_0$. These conditions can be treated as matching conditions for solutions of the linear equations at $T<T_0$ and $T>T_0$. As a result we may obtain the following system of equations instead of the system~(\ref{sys:simpl}) for monotonic in $\xi$ functions. It is convenient for the investigation of the stationary traveling wave and its linear stability: for $T_{+}>T_0$ \begin{equation} \partial_t T_{+}=v \partial_{\xi} T_{+} +\partial_{\xi}^2 T_{+}+\omega_0 c_{+},~~\partial_t c_{+}=v \partial_{\xi} c_{+}-\omega_0 c_{+}, \label{moving_1} \end{equation} and for $T_{-}<T_0$ \begin{equation} \partial_t T_{-}=v \partial_{\xi} T_{-} +\partial_{\xi}^2 T_{-},~~c_{-}=1. \label{moving_2} \end{equation} Here \begin{equation} f(\xi,t)=\left\{ \begin{array}{lcr} f_{-}(\xi,t)&\mbox{~~for~~}& \xi>\xi_f(t),\\ f_{+}(\xi,t)&\mbox{~~for~~}& \xi<\xi_f(t), \end{array} \right. \label{moving_3} \end{equation} where $f$ represents $T$ or $c$, and $\xi_f(t)$ is position of the front, where $T=T_0$. The matching conditions read as: \begin{eqnarray} T_{+}(\xi_f(t),t)&=&T_{-}(\xi_f(t),t)=T_0;\nonumber\\ \partial_\xi T_{+}(\xi_f(t),t)&=&\partial_\xi T_{-}(\xi_f(t),t);\nonumber\\ \quad c_{+}(\xi_f(t),t)&=&1. \label{match} \end{eqnarray} The stationary traveling wave obeys the following boundary conditions \begin{equation} \xi\rightarrow\infty:~T=0,~c=1;~~\xi\rightarrow -\infty:~\partial_x T=0,~~c=0, \label{BC} \end{equation} and the conditions: \begin{equation} \partial_t c=\partial_t T=\xi_f(t)=0 \, . \label{stat} \end{equation} The latter equality originates from a freedom due to the translation invariance of the problem, and hence from possibility to set the front at an arbitrary point of the $\xi$-space. The system~(\ref{moving_1})-(\ref{stat}) presents nonlinear eigen-value problem, with $v$ being the eigen-value. There is unique solution of this problem: \begin{eqnarray} &&~~v=1\nonumber\\ &\xi>0:&~~c=1,~~T=T_{-}=T_0e^{-\xi},\nonumber\\ &\xi<0:&~~c=c_{+}=e^{\omega_0\xi},~~T=T_{+}=1-\frac{1}{\omega_0+1}e^{\omega_0\xi}\, . \label{sys:simpl_solution} \end{eqnarray} As a result we will set $v=1$ as a constant velocity of the moving frame for the equations~(\ref{moving_1}) and~(\ref{moving_2}). We have for the traveling burning wave velocity in the initial physical units: $$ v=\sqrt{\frac{\kappa\omega}{C}\,\left(\frac{T_f}{T_0}-1\right)}\, . $$

We have presented here a simple analytically-solvable model for flame propagation which exhibits a different behavior depending on the dimensionless parameter {\sl Ze}: the effective Zeldovich number expressed through our dimensionless parameter $\omega_0$. The theory is compared with numerical simulations, and a good agreement is observed. When {\sl Ze}~$<6$ both analytic and numerical solutions exhibit a constant-velocity front. Our analytical theory shows that the traveling flame front becomes unstable at {\sl Ze}~$>6$, whereas our numerical simulations show that the front is destroyed and passes to jerk--like pulsations. We believe that, in the case of pulsating instability in a real flame, the characteristic behavior will be the same as in this model, but the critical {\sl Ze} can be different. Therefore, the proposed model can mimic the behavior of numerical methods, and serve as a test for them. Our present work cannot say anything about the existence of such pulsations in real white dwarfs. However, if such a pulsating jerk-like regime of slow flame propagation indeed takes place in white dwarfs, then it could be able to trigger the transition to detonation. We gratefully acknowledge extensive discussions with S.~I.~Blinnikov and B.~Meerson. The work is supported partly by grants RFBR 10-02-00249-a, RFBR 11-02-00441-a, ``Dynasty'' foundation, SCOPES project No.~IZ73Z0-128180/1, by Federal Programm ``Scientific and pedagogical specialists of innovation Russia'' contract number 02.740.11.0250, and by the contract No.~02.740.11.5158 of the Ministry of Education and Science of the Russian Federation.

1012.5982

1012

1012.2863.txt

{ We present the results of the spectral analysis of the public data of 438 Gamma Ray Bursts (GRBs) detected by the \fe\ Gamma ray Burst Monitor (GBM) up to March 2010. For 432 bursts we could fit the time integrated spectrum. In 318 cases we can reliably constrain the peak energy \epo\ of their $\nu F_{\nu}$ spectrum by analyzing their time--integrated spectrum between 8 keV and 35 MeV. 80\% of these spectra are fitted by a power--law with an exponential cutoff, and the remaining with the Band function. Among these 318 GRBs, 274 and 44 belong to the long and short GRB class, respectively. Long GRBs have a typical peak energy \epo$\sim$160 keV and low energy spectral index $\alpha\sim -0.92$. Short GRBs have harder peak energy (\epo$\sim$490 keV) and harder low energy spectral index ($\alpha\sim -0.50$) than long bursts. For each \fe\ GRB we analyzed also the spectrum corresponding to the peak flux of the burst. On average, the peak spectrum has harder low energy spectral index but similar \epo\ than the corresponding time--integrated spectrum for the same burst. The spectral parameters derived in our analysis of \fe/GBM bursts are globally consistent with those reported in the GRB Cicular Network (GCN) archive after December 2008, while we found systematic differences, concerning the low energy power law index, for earlier bursts. }

Our current knowledge of the spectral properties of the prompt emission in GRBs mainly relies on the data collected in almost 10 years by the Burst And Transient Source Experiment (BATSE) onboard the {\it Compton Gamma--Ray Observatory (CGRO)}. \ba\ allowed to characterize the spectrum of the population of short and long GRBs over a large energy range, from 20 keV to 1--2 MeV. The analysis of such data revealed some important results about the spectral properties of these GRBs. The prompt spectra, integrated over the GRB duration (i.e. time integrated spectra), can be typically well described by a curved function showing a peak -- in a $\nu F_\nu$ representation -- at a typical energy \epo\ of a few hundreds of keV but whose distribution spans nearly three orders of magnitude. Large dispersions characterise also the distributions of the low--and high--energy photon indices, whose characteristic values are $\alpha\sim -1$ and $\beta\sim -2.3$, respectively (Band et al. 1993; Ghirlanda et al. 2002; Kaneko et al. 2006). Similar results are obtained by considering the time resolved spectral analysis of flux/fluence limited samples of bright \ba\ bursts (Preece et al. 1998; Preece et al. 2000; Kaneko et al. 2006). The BATSE data also suggested the existence of two different classes of GRBs (long and short), based on both temporal and spectral features. Evidence of a spectral diversity between long and short bursts comes from their different hardness--ratios (HR) (Kouveliotou et a. 1993). The larger HR of short bursts might be ascribed to a larger \epo. Nava et al. (2008) and Ghirlanda et al. (2009 -- G09 hereafter) showed that \epo\ correlates both with the fluence and the peak flux. Although short and long bursts follow the very same \epo--peak flux relation, they obey different (parallel) \epo--fluence relations. This implies, obviously, that the distributions of the ratio \epo/fluence is different for the two burst classes. Recently, Goldstein, Preece \& Briggs (2010) proposed this ratio as discriminator between short and long GRBs. Due to the relation between \epo\ and the bolometric fluence and peak flux, a direct comparison between the \epo\ distributions of the two different burst classes must take into account the different fluence/peak flux selection criteria. G09 analyzed and compared samples of short and long BATSE bursts selected with similar peak flux limits. They found that the peak energy distributions of the two classes are similar, while the most significant difference concerns the low--energy power--law indices, with short bursts having typically a harder $\alpha\sim -0.4$. These global spectral properties of GRBs have also been confirmed by other satellites ({\it Beppo}SAX, Hete--II and Swift) (Guidorzi et al. 2010; Sakamoto et al. 2005; Butler et al. 2007). However, the detectors on board these satellites have different sensitivities with respect to \ba\ and cover a narrower and different energy range. For instance, the relatively narrow energy range (15--150 keV) of \sw/BAT does not allow to constrain the spectral peak \epo\ for most of the detected bursts (Cabrera et al. 2007; Butler et al. 2007). Spectral studies of the prompt emission of GRBs require a wide energy range, possibly extending from few tens of keV to the MeV energy range. This allows to measure the curvature of the GRB spectrum and to constrain its peak energy, as well as its low and high energy spectral slopes. The \fe\ satellite, launched in June 2008, represents a powerful opportunity to shed light on the origin of the GRB prompt emission thanks to its two instruments: the Large Area Telescope (LAT) and the Gamma--ray Burst Monitor (GBM). LAT detected in about 2 years very high energy emission ($>$ 100 MeV) from 19 GRBs. This emission component shares some common features with that already found in few bursts by EGRET (Energetic Gamma--Ray Experiment Telescope) on board the {\it CGRO} satellite. In particular, high energy $\sim$GeV flux is still observed when the softer energy emission (in the sub--MeV domain) is ceased, and often its onset lags the sub--MeV component (e.g. Ghisellini et al. 2010, Omodei et al. 2009). The other instrument on board \fe\ is the GBM (Meegan et al. 2009) which is similar to \ba\ and, despite its slightly worse sensitivity, allows to study the GRB spectrum over an unprecedented wide energy range, from 8 keV to 40 MeV. This is achieved by its twelve NaI detectors giving a good spectral resolution between $\sim$8 keV and $\sim$1 MeV and two BGO detectors which extend the energy range up to $\sim$40 MeV. Similarly to \ba, the NaI detectors guarantee full--sky coverage, but their smaller geometric area (16 times lower than that of the LADs of \ba) implies a lower sensitivity. On the other hand, the presence of the BGO detectors allows for the first time to extend the energy range for the study of the spectrum of the prompt emission to tens of MeV thus accessing an energy range poorly explored with the CGRO instruments. 438 events, classified as GRBs\footnote{http://heasarc.gsfc.nasa.gov/W3Browse/fermi/fermigbrst.html}, triggered the GBM until the end of March 2010. {\it With this large sample of bursts we can perform, for the first time, a robust statistical study of the spectral properties of \fe/GBM bursts}. The main aims of this paper are: (a) to present the results of the spectral analysis of 438 \fe\ bursts, (b) to show the distribution of their spectral parameters, (c) to compare the spectral properties of \fe\ short and long GRBs and (d) to compare the spectra integrated over the burst duration (time--integrated spectra, hereafter) with the spectra of the most intense phase of the burst, i.e. its peak flux (peak spectra, hereafter). Preliminary results of the spectral analysis of \fe\ GRBs performed by the GBM team have been distributed to the community through the Galactic Coordinates Network (GCN) Circulars. These amount to 228 GRBs (until March 2010), whereof 167 have a well constrained \epo. On going spectral calibrations of the GBM detectors make the results published in the GCN ``preliminary'', especially for the first bursts detected by the GBM. The GBM team continuously provides, together with the public data of detected GRBs, more updated detector response files. Few months ago the software and the new response files have been made public so that a systematic and reliable analysis of the spectra of \fe/GBM bursts is now possible. We will compare the results of our spectral analysis with those published in the GCN to search for possible systematic effects in the GCN results. The paper is organized as follows: in \S 2 and \S 3 we describe the sample of \fe\ GRBs and the procedure adopted to extract and analyze their spectral data, respectively. In \S 4 we present the spectral results and build their distributions, also considering short and long GRBs separately. In \S 4 we also compare the time--integrated spectra and the peak spectra for the analyzed bursts. We summarize our results in \S 5.

We analyzed the spectra of all GRBs detected by the \fe\ Gamma-ray Burst Monitor (GBM) between 14 July 2008 and 30 March 2010. These are 438 GRBs and for 432 of them we have all the needed data to perform the spectral analysis. The time--integrated spectrum is best fitted with a power law model (110 spectra -- reported in Tab. \ref{tab11}) or a curved model (323 spectra -- reported in Tab. \ref{tab1}) which is either the Band model (65 spectra) or a cutoff--power law (CPL) model (258 spectra). Among the 432 GRBs for which we could analyze the spectrum, we identify 73 short and 359 long bursts, respectively. Their $LogN-LogF$ is similar (Fig. \ref{fg0}) and its high--fluence tail is consistent with a powerlaw with slope -3/2. The 73\% of the bursts detected by the GBM up to March 2010 could be fitted with a curved model (Band or CPL, with a prevalence of the latter model) and in the majority (318 out of 323) of these cases we could constrain the spectral parameters and in particular the peak energy \epo\ of the $\nu F_{\nu}$ spectrum. This is possible thanks to the large energy range of the GBM spectra extending from 8 keV to $\sim$35 MeV. This is the sample we considered for the characterization of the spectral parameters of the time--integrated spectra of \fe\ GRBs. Within this sample there are 44 short and 274 long GRBs. The comparison of their spectral properties shows that short GRBs have higher \epo\ than long events (Fig. \ref{fg1}) and a slightly harder low energy spectral index $\alpha$ (Fig. \ref{fg2}). The finding that short \fe\ GRBs have harder peak energy than long events seems opposite to what found from the comparison of short and long GRBs detected by BATSE (Ghirlanda et al. 2009). However, the \fe\ short GRBs have also larger peak fluxes than long events. A more detailed comparison between long and short GRBs detected by {\it Fermi}/GBM and BATSE is presented in Nava et al. 2011. A second major part of the present work was aimed to characterize the spectra of the peak of each GRB. Through time resolved spectroscopy we isolated and analyzed the spectrum corresponding to the peak of the flux light curve of each burst. The results are reported in Tab. \ref{tab3}. By comparing the peak spectrum and the time--integrated spectrum of individual GRBs we find that the peak spectra have similar \epo\ of the time integrated spectra but harder low energy spectral index $\alpha$ (Fig. \ref{fg5}). Finally we compared the results of our spectral analysis with those reported in the GCN circulars. We found that, due the still not fully completed calibrations of the GBM detectors, the GCN results of bursts comprised between July and December 2008 are affected by a systematic overestimate of the hardness of the GRB spectrum at low energies (i.e. the spectral parameter $\alpha$). This systematic bias does not affect \epo\ and is not present in our results which are obtained with the most recent releases of the GBM response files.

1012.2863

1012

1012.4330_arXiv.txt

We report on the analysis of all 65 pointed {\it Rossi X-ray Timing Explorer} observations of the recently discovered soft X-ray transient MAXI J1659--152 (initially referred to as GRB 100925A). The source was studied in terms of its evolution through the hardness-intensity diagram (HID) as well as its X-ray variability properties. MAXI J1659--152 traced out an anti-clockwise loop in the HID, which is commonly seen in transient low-mass X-ray binaries. The variability properties of the source, in particular the detection of type-B and type-C low-frequency quasi-periodic oscillations, and the way they evolve along the HID track, indicate that MAXI J1659--152 is a black hole candidate. The spectral and variability properties of MAXI J1659--152 imply that the source was observed in the hard and soft intermediate states during the {\it RXTE} observations, with several transitions between these two states.

Black hole X-ray binaries have been studied since the early 1970s. It is now well established that the X-ray spectral and variability properties of these sources are strongly correlated, which is most clearly seen in the transient black hole X-ray binaries (BHTs). While there is a great variety in the observed outburst behavior of BHTs (even for single sources), their outbursts typically proceed along 'q'-shaped tracks in hardness-intensity diagrams (HIDs; see, e.g., \citealt{Homan05} and \citealt{Dunn10}), which are traced out in an anti-clockwise manner; we note that similar tracks are also traced out by neutron star transients \citep[e.g.][]{Tudose09,Linares09}. The various branches of these tracks correspond to distinct spectral states. Not all sources show all possible states and the time they spend in each state can differ significantly. There are various conventions for describing the spectral (and variability) states of BHTs; see, e.g., reviews by \citet{Homan05}, \citet{Remillard06}. In this paper we follow the convention used in \citet{Belloni10b}, which is based on the work by \citet{Belloni05} on GX 339-4. This source is often used as a template for the outburst evolution of BHTs, owing to its well defined q-shaped HID tracks, and the fact that it shows behavior that is common to many other systems. When GX 339-4 goes into outburst, its intensity is low and spectrum is hard (low-hard state or LHS). As the intensity increases, the spectrum remains hard, until the source makes a transition to the intermediate state (IMS), where the hard color starts to decrease at a rather constant intensity. The IMS can be divided into a hard and a soft IMS (HIMS and SIMS, respectively) depending on the spectral and variability characteristics observed. The transition from the LHS is always first to the HIMS. GX 339-4 often shows several transitions between HIMS and SIMS until the hardness decreases even further and the source reaches the so called high-soft state (HSS), where subsequently the hardness remains approximately constant as the intensity eventually decreases. At some point during this decrease hardness increases again and the source transits from the HSS via the IMS back to the LHS, and returns to quiescence. A BHT exhibits various types of quasi-periodic oscillations (QPOs) and broad-band noise components, whose properties are strongly correlated with spectral state \citep[see, e.g.,][]{Klein07}. The power spectrum of the LHS is characterized by strong broadband variability (0.01-100 Hz fractional rms amplitude up to 50\%). During the HIMS so-called type-C QPOs are observed \citep[see][for QPO type definitions]{Wijnands99,Remillard02,Casella05}, often with strong harmonic content and accompanied by strong broadband variability (fractional rms up to 30\%). In the SIMS various types of variability are observed: power spectra with type-B QPOs (also with strong harmonic content) or (weaker) type-A QPOs, but also power spectra with weak peaked noise and/or QPO features that have been poorly characterized. All SIMS power spectra have in common that their broadband variability is (considerably) weaker than in the LHS and HIMS. Type-C QPOs are typically observed between 0.1--10 Hz, whereas type A and B QPOs are generally confined to the 4--8 Hz range. In the HSS variability reaches its minimum strength (a few percent rms); sometimes very weak QPOs above 10 Hz are seen \citep{Homan01}. In this paper we present a study of {\it Rossi X-ray Timing Explorer (RXTE)} observations of MAXI J1659--152 (henceforth J1659). This source was discovered with the {\it Swift} Burst Alert Telescope \citep[BAT;][]{Barthelmy05} on September 25, 2010 and was initially thought to be a gamma-ray burst \citep[GRB 100925A, see][]{Mangano10}. Later, it was suggested by \cite{Kann10} to be a new Galactic X-ray transient due to its persistent X-ray emission; this was confirmed by \cite{Negoro10} with {\it MAXI} observations. J1659 has also been detected in the radio \citep{Vanderhorst10}, optical \citep{Jelinek10} and sub-mm bands \citep{deugartepostigo10}. The aim of this letter is to discuss the results of the aperiodic timing analysis and color evolution from {\it RXTE} observations, based on which we conclude that J1659 is a black hole candidate \citep{Kalamkar10}.

We analyzed all the {\it RXTE} observations of MAXI~J1659--152 taken during its 2010 outburst. Spectral states were identified using the HID in combination with the aperiodic variability properties. In the HID J1659 traced out an anti-clockwise loop, similar to those seen in black hole and neutron-star transients (see \S\ref{intro}). The types of power spectral features we identified (in particular type--B and type--C QPOs) and the way they evolve along the HID track, indicates that J1659 is a black hole candidate. The spectral decomposition of the softest spectrum (i.e.\ diskbb + power-law, without the need for second thermal component; \citealt{Lin07}), the radio loudness in the early part of the outburst \citep{Paragi10}, the flip-flops seen in two observations, and the sudden ($\sim$day) drop in the rms (i.e. PDS3) at low-hardness, are all characteristic of this type of systems \citep{Dunn10,Miyamoto91,Fender09}, strengthening our identification of J1659 as a black hole candidate. Based on spectral and variability properties we find that J1659 moved from a HIMS to a SIMS and back again (at lower intensity), without reaching the HSS (or the soft or thermal dominant state as defined by \citealt{Remillard06}). Our {\it RXTE} observations started after the source had left the LHS (as indicated by {\it MAXI} observations), and ended before it had returned to that state. The observations in the softest part of the SIMS had weak variability and revealed two types of power spectra, one consistent with a power-law and one with an additional bump around 7--8 Hz. We note that similar groups of observations with weak variability were observed in the SIMS of GX 339--4 by \citet{Belloni05}. \citet{Fender09} found that radio flares and subsequent quenching of the radio flux often occur in a time interval of a few days before and after the time of sudden drops in the rapid X-ray variability (identified as a distinct zone in their rms-hardness diagrams). In J1659 a similar drop occurred on MJD 55481, a few days after the quenching of the radio flux as reported by \citet{Vanderhorst10a}, consistent with other BHTs. Optical observations constraining the mass of the compact object are required to confirm the black hole nature of J1659. From dipping episodes in the X-ray light curves, the orbital period period has been proposed to be 2.4--2.5 hour \citep{Kuulkers10,Belloni10a}. This makes the source very interesting as, if confirmed, this would make it the black hole binary with the shortest known orbital period. \textbf

1012.4330

1012

1012.1515_arXiv.txt

Studies of a class of infinite one dimensional self-gravitating systems have highlighted that, on the one hand, the spatial clustering which develops may have scale invariant (fractal) properties, and, on the other, that they display ``self-similar" properties in their temporal evolution. The relevance of these results to three dimensional cosmological simulations has remained unclear. We show here that the measured exponents characterizing the scale-invariant non-linear clustering are in excellent agreement with those derived from an appropriately generalized ``stable-clustering" hypothesis. Further an analysis in terms of ``halos" selected with a friend-of-friend algorithm reveals that such structures are, statistically, virialized across the range of scales corresponding to scale-invariance. Thus the strongly non-linear clustering in these models is accurately described as a virialized fractal structure, very much in line with the ``clustering hierarchy" which Peebles originally envisaged qualitatively as associated with stable clustering. If transposed to three dimensions these results would imply, notably, that cold dark matter halos (or even subhalos) are 1) not well modeled as smooth objects, and 2) that the supposed ``universality" of their profiles is, like apparent smoothness, an artefact of poor numerical resolution.

The much acclaimed successes of the $\Lambda$CDM cosmology in matching many observations concern essentially its homogeneous limit and the linear regime in which perturbations to homogeneity are small. The success of the model in accounting for the numerous observations which probe the non-linear regime, where density fluctuations are large, is much more uncertain. An important consideration in this respect is the great difficulty of calculating the model's predictions in this regime. Even in the idealized limit in which clustering arises from gravity alone, predictions reside solely on numerical simulations. The latter have, despite impressive increases in their size and sophistication, still quite limited spatial resolution. Further, essentially because there are no non-trivial analytical benchmarks which they can be tested against, it remains unclear to what extent their results are conditioned by these resolution limits. In this paper we explore what might be learnt about the nature of non-linear gravitational clustering from the study of a class of one dimensional (1D) models, which have the interest of offering, even with modest computer resources, very much greater spatial resolution and exact numerical integration. An early study in this spirit is that of \cite{melott_prl_1982,melott_1d_1983} which used such a 1D model to explore clustering in hot dark matter cosmologies. \cite{rouet_etal} derived and studied a slightly different model to the most naive 1D version of clustering in an Einstein de Sitter (EdS) universe considered by \cite{yano+gouda, aurell+fanelli_2002a, aurell_etal_2003}. \cite{valageasOSC_2} and \cite{agmjfs_pre2009} have studied the version without expansion, while Miller et al. \citep{miller+rouet_2002, miller+rouet_2006, miller_etal_2007, miller+rouet_2010a} have reported results on all three of these models. Despite these studies it remains unclear, however, whether a good, and really useful, analogy can be made with 3D cosmological simulations. Generalizing the results of \cite{ yano+gouda} and \cite{agmjfs_pre2009}, we confirm here the very strong qualitative similarities in the temporal development of clustering to that in the 3D case. Further we show the scale invariant properties of the clustering in real space, emphasized and analyzed by \cite{ miller+rouet_2002, miller+rouet_2006, miller_etal_2007}, can in fact be well explained within an analytical framework similar to one proposed long ago for the 3D case \citep{peebles_1974, davis+peebles_1977, peebles}. We derive the appropriate 1D version of this generalized ``stable clustering" prediction for the correlation exponents, and show it to provide a very good approximation to the numerical results. We also show that, in a precise sense, the strongly non-linear clustering can be properly described as a {\it virialized} fractal hierarchy, in line with what was originally anticipated qualitatively in three dimensions. We argue finally for the extrapolation of these conclusions about the qualitative nature of clustering to the relevant 3D case, and consider its implications, in particular for what concerns the nature and properties of ``halos" in cold dark matter cosmologies. \section {1D versions of cosmological N-body simulations} Dissipationless cosmological N-body simulations (for a review see, e.g., \cite{bagla_review}) solve numerically the equations \begin{equation} \label{3d-equations} \frac{d^2 {\bf x}_i}{dt^2} + 2H \frac{d{\bf x}_i}{dt} = -\frac{Gm}{a^3}{\sum}^J \frac{{\bf x}_i - {\bf x}_j}{\vert {\bf x}_i - {\bf x}_j \vert^3} \,, \end{equation} where ${\bf x}_i$ are the comoving particle coordinates, $a(t)$ is the appropriate scale factor for the cosmology considered, with Hubble constant $H(t)={\dot a}/{a}$. The superscript `J' in the sum, which runs in practice over the infinite system constituted by periodic copies of a cube containing $N$ particles, indicates that the (badly defined) contribution of the mean mass density (which sources the Hubble expansion) is subtracted. We consider here simply the 1D system obtained by replacing the 3D Newtonian force term by the analogous 1D expression, derived starting from the 1D Poisson equation which gives a force between particles in one dimension (or, equivalently, infinite parallel sheets embedded in three dimensions) independent of separation. We thus consider the equations \begin{equation} \label{1d-equation} \frac{d^2 x_i}{dt^2} + 2H \frac{d x_i}{dt}= -\frac{g}{a^3} \lim_{\mu \rightarrow 0} \sum_{j\neq i} \textrm{sgn}(x_i - x_j) e^{-\mu \vert { x}_i - {x}_j \vert}, \end{equation} where the limiting procedure used in the sum is simply a convenient way to explicit the subtraction of the background \citep{kiessling}, and $g$ is the coupling constant (with $g \equiv 2\pi \Sigma G$ for sheets of surface mass density $\Sigma$). Initial conditions are generated, just as in 3D cosmological simulations, by applying appropriate small perturbations to particles initially on an infinite perfect lattice. If the displacement from its initial lattice site of particle $i$ is $u_i$ it can be shown rigorously \citep{agmjfs_pre2009} that the net force on the particle, until it crosses another one, is exactly proportional to $u_i -\langle u \rangle$, where $\langle u \rangle$ is the average displacement. The equations of motion up to the time at which particles cross become \begin{equation} \label{1d-equation-u} \ddot{u}_i + 2H \dot{u}_i= \frac{2gn_0}{a^3}u_i \end{equation} where $n_0$ is the mean particle density (and we take $\langle u \rangle=0$). When the {\it particles} in our model cross one another the force changes simply by $\pm 2g$. However, since in 1D a crossing of particles is equivalent, up to exchange of labels, to an elastic ``collision" between them, one can instead consider (as we are not interested in the labels of the particles) a system of particles which bounce elastically and follow Eq ~(\ref{1d-equation-u}) {\it at all times} between ``collisions". Finally, by a transformation of the time coordinate to $\tau= \int dt/a^{3/2}$, and appropriate choice of units of $\tau$, we can rewrite these equations as \begin{equation} \label{1d-equation-u-final} \frac{d^2 {u}_i}{d\tau^2} + \Gamma \frac{d { u}_i}{d\tau} = u_i \,, \end{equation} i.e., as those of a set of damped inverted harmonic oscillators. We note that these equations coincide simply with those for displacements of {\it fluid} elements in the Zeldovich approximation (see, e.g. \cite{buchert2}), which is in fact exact up to ``shell crossing" for 1D perturbations. The same system can thus be derived as a simple analytical continuation of this approximation at shell crossing \citep{yano+gouda, aurell_etal_2003}. For a generic cosmology $\Gamma$ in Eq.~(\ref{1d-equation-u-final}) is a non-trivial function of $\tau$, but in the specific case of an EdS cosmology (for which $a \propto t^{2/3}$) it is a constant, $\Gamma=1/\sqrt{6}$ (and $\tau \propto \ln a$). The model derived and studied by \cite{rouet_etal} corresponds instead to the case $\Gamma=1/\sqrt{2}$, while the case $\Gamma=0$ (i.e. static limit) has also been studied by several authors (see references given above). We note that, despite the fact that we obtained the value $\Gamma=1/\sqrt{6}$ above, this is not necessarily {\it the} more appropriate value to consider for our study: in the derivation just given the 3D Hubble law has been imposed {\it by hand}. This has been done because the 1D Hubble law (which would arise if one follows fully the analogy with the 3D derivation starting from physical coordinates) gives a completely different qualitative behavior --- collapse in a finite time --- which has no relevance to 3D cosmology. Thus in order to ``mimic" EdS cosmology, one could equally take the functional form of the corresponding Hubble law, but leave its normalization as a free parameter. This leads to the Eqs.~(\ref{1d-equation-u-final}) with $\Gamma$ an undetermined constant. These equations then simply define a simple class of toy models, which may be useful for {\it qualitative} understanding of the 3D problem. Given initial displacements and velocities, Eqs.~(\ref{1d-equation-u-final}) can trivially be integrated exactly between particle collisions when $\Gamma$ is a constant. To determine the time of the next collision (and which pair of particles it involves) requires then only the solution of algebraic equations. The numerical integration can therefore be performed ``exactly", i.e., up to machine precision. As described by \cite{noullez_etal} the integration can be sped up optimally using a ``heap" structure. All the results presented here are for systems with $N=10^5$ particles (and periodic boundary conditions). In making the analogy with 3D simulations, which aim to reproduce the {\it collisionless} limit of gravitational clustering and include a smoothing of the singularity in the 3D Newtonian force for this reason, it is to be noted that, on the time scales we will consider, these 1D systems, without force smoothing, can be expected to represent extremely well this limit. Studies of collisional relaxation in finite 1D self-gravitating systems (see \cite{mjtw_relaxation_2010}, and references therein) have shown that it is very suppressed compared to that in 3D systems, a typical relaxation time for a virialized system of number density $n$ being $\sim (10^4 - 10^6) N/\sqrt{gn}$. Using this estimate in the simulations reported below, it is simple to show that even the relaxation time for the smallest and densest clusters is much longer than the duration of the simulation. The reason for this relative suppression of this relaxation is that, in one dimension, there is no analogy to 3D 2-body relaxation: the ``collisions" of particles we have discussed above do not contribute to the usual collision term in the Boltzmann equation. Indeed it is for this reason that particle crossings and ``collisions" are equivalent.

{\it In three dimensions} the stable clustering hypothesis has been used widely as a reference point, and indeed much used phenomenological models such as the formalism proposed by \cite{peacock} incorporate it into the modeling of the non-linear power spectrum obtained from numerical simulations. However, in parallel, results of major simulations of power law initial conditions (e.g. \cite{efstathiou_88, smith}) led to the conclusion that, although the measured exponents in the correlation function showed a behaviour roughly consistent with its predictions (albeit with some significant deviations \citep{smith}), there was no evidence for the ``clustering hierarchy" which Peebles had argued would be associated with it (see, e.g. \cite{peebles}). \cite{efstathiou_88}, notably, explicitly excluded such a ``fractal" description and found evidence instead for the validity of a description in terms of ``smooth non-linear clumps". These latter are the precursors of the ``halos" of halo models, which have become the standard phenomenological description of the matter distribution in cold dark matter cosmologies (for a review, see, e.g. \cite{halo}). The matter distribution is then approximated as a collection of isolated (and thus virialized) spherical structures with smooth density profiles. Further, on the basis of extensive numerical study following that of \cite{navarro2}, the latter are widely believed to be characterized well by ``universal" exponents (i.e. independent of initial conditions and cosmology). \begin{figure*} \begin{center} \includegraphics[angle=-90, width=2\columnwidth]{fig3.eps} \caption{ Density field obtained at $t=\log a=14$ starting from an initial condition with $n=2$, in the ``EdS" model with $\Gamma=1/\sqrt{6}$. The first panel shows the whole box (i.e. length units are such that $L=N=10^5$), while each subsequent panel shows a spatial ``zoom" on a region of size {\it one tenth} that shown in the previous panel. The resolution has been increased (i.e. the bin size has been decreased) in the last three plots in order to reveal the clumpy nature of the distribution at these scales } \label{Zoom_density_fractal} \end{center} \end{figure*} {\it In one dimension } the strongly non-linear regime is truly scale-invariant (for power law initial conditions) in a range which grows monotonically in time. The associated distribution of matter is {\it intrinsically} lumpy or grainy down to the lower cut-off scale $x_{min}$: indeed the very meaning of such {\it scale invariance} is that there are no characteristic scales available to define smoothness. To illustrate this we show in Fig. ~\ref{Zoom_density_fractal} the spatial density at various levels of detail in a typical evolved configuration. Such distributions are most naturally described with the instruments of fractal (or, more generally, multi-fractal) analysis developed for this purpose in condensed matter and statistical physics (see e.g. \cite{book}). As we have just illustrated in the previous section, halo type descriptions may, of course, also be employed to describe them. The halos so defined are, however, very different to those described by 3D halo models: these 1D halos are not smooth. Further have they no intrinsic size themselves, but are defined only with reference to an arbitrary chosen scale. The study of the virial ratios we have presented indicates, however, that such halos can be considered as entities with a dynamical relevance, as they show a clear tendency to have a virial ratio of order unity (which is the behaviour of an isolated structure). The clustering in the non-linear regime can thus be considered as a concrete realization of the qualitative picture of a ``clustering hierarchy" originally envisaged by Peebles (e.g. \cite{peebles}) as resulting from stable clustering. The stable clustering hypothesis we have described above, however, is actually subtly different from the original one: we assumed only that stable clustering applies below the scale $x_{min}$ marking the lower cut-off to the scale invariance, and not necessarily to the strongly non-linear regime as a whole. Thus we assumed only that stable clustering applies at an ultraviolet scale {\it fixed by the resolution of the simulation} (or, physically, by the scale at which the very first structures form). The ``statistical virialization" we have observed using the halo analysis, on the other hand, applies at scales above $x_{min}$ and across the range of the scale invariant clustering. There are clearly two possible conclusions one can draw from this analysis: \begin{itemize} \item A. These 1D models produce non-linear clustering which is qualitatively different in its nature to that in 3D, or \item B. The spatial resolution in 3D simulations up to now has been too limited to reveal the nature of clustering in cold dark matter cosmologies, which is correctly reflected (qualitatively) in the 1D simulations. \end{itemize} We believe that, despite the impressive computational size and sophistication of 3D cosmological simulations, conclusion B may well be the correct one. The very largest modern studies in a cosmological volume access roughly two decades in scale in the non-linear regime\footnote{ For example, the ``millennium" simulation \citep{springel_05} of the $\Lambda CDM$ cosmology, has, at redshift zero, $\xi=1$ at about $2$h$^{-1}$Mpc, an initial mean interparticle $\ell \approx 250$h$^{-1}$ kpc and a force smoothing $\epsilon=5$h$^{-1}$ kpc. It is usually supposed that it is the latter which sets the lower bound on resolution, but we note that it has been shown by several studies (see \cite{melott_etal_1997, romeo_2008, mjbmtb_2009}) that, at the very least, precision may be compromised below the scale $\ell$.} while reference studies in the literature of power law initial conditions in EdS cosmology \citep{efstathiou_88, smith} measure the crucial power-law behaviour in the correlation function (or the PS) over at most one decade. If we were to perform our 1D simulations at comparable resolution to large cosmological simulations like \cite{smith}, we would certainly have great difficulty in establishing the scale-invariant nature of the strongly non-linear clustering arising from power law initial conditions. Although halos defined exactly as in three dimensions might look clumpy, an approximately smooth profile could be determined for them if they were averaged (as they can be in three dimensions when spherical symmetry is assumed). Higher resolution 3D simulations of smaller regions have shown over the last decade that there is in fact much more substructure (``subhalos") inside halos than was originally anticipated (see, e.g., \cite{moore_etal_1999, diemand_etal_2005, goerdt_etal_2007}), and have even more recently described several levels of such substructure (``subhalos of subhalos", see e.g. \cite{diemand_etal_nature_2008, springel_etal_aquarius_2008, stadel_etal_2009}). \cite{diemand_etal_nature_2008} even use the term ``fractal" to describe (qualitatively) the real space structures, while \cite{zemp_etal_2009} describe the structure of halos in phase space as ``intrinsically grainy". We note that other authors (see, e.g., \cite{valageas_1999, gaite_2007}) have previously argued for similar conclusions on the basis of analyses of 3D simulations. Let us consider nevertheless one possible consideration in favour of (the more conservative) conclusion A. In the expanding (i.e. damped) 1D models, the stable clustering prediction (\ref{gamma-sc-1D}) fits the measured exponents extremely well. Early 3D studies for EdS cosmologies (e.g. \cite{efstathiou_88}) measured exponents roughly consistent with the stable clustering prediction, but later studies (e.g. \cite{smith}) have found significant disagreement. This disagreement is attributed to physical mechanisms which cause the fundamental assumption of stability to be violated --- by the evident fact that {\it there are} interactions between ``halos", which can even lead to their merging into single structures. We have noted that in one dimension tidal forces vanish, and structures can interact only when they actually physically cross one another. While merging may occur, it may be that it is a less efficient process than in three dimensions. Thus the excellent agreement in the 1D models compared to EdS may perhaps be attributed to the fact that these models probably represent poorly the role of such physical effects. The essential question, however, is not whether these effects play a role and can lead to deviations from stable clustering, but whether such effects can {\it qualitatively} change the nature of clustering, destroying scale invariance by smoothing out the distribution {\it on a scale related to the upper cut-off} to scale invariance. Our study of the case $\Gamma=0$ suggests that the answer is negative. The prediction of stable clustering does not work in this case, and like in three dimensions, one obtains a small value of the exponent which does not sensibly depend on $n$. The physical reasons why the exponent is close to, but different to, the stable clustering prediction are a priori the ones just cited. The analysis of Miller et al. of this case, which we have rechecked and confirm, finds nevertheless that the distribution is scale invariant. Further, as we have mentioned, the lower cut-off $x_{min}$ remains constant as in the stable clustering hypothesis, of order the initial lattice spacing (and unrelated to the upper cut-off). These results on 1D models suggest directions for 3D investigations which might establish definitively the correctness of conclusion B. We note, for example, that the 1D models lead one to expect that the exponents derived phenomenologically to characterize the highly non-linear density field inside smoothed halos (i.e. the ``inner slope" of halos) should be closely related to the exponent $\gamma$ determined from the correlation function. Indeed --- in the approximation of a simple fractal behavior in the strongly non-linear regime, which the spectrum of multi-fractal exponents measured in \cite{miller+rouet_2010a} suggests should be quite good --- the mean density about the centre of such halos will decrease just as about any random point, i.e., with the {\it same} exponent $\gamma$. Despite the contradiction with the widely claimed ``universality" of such exponents in halos profiles, such a hypothesis cannot currently be ruled out, as the determination of such exponents is beset by numerical difficulties (arising again from the limited resolution of numerical simulations). In a study of halo profiles obtained from power law initial conditions \cite{knollmann_etal_2008} show explicitly that the results for the halo exponents depend greatly on what numerical fitting procedure is adopted. While one procedure gives ``universality" (i.e. exponents independent of $n$), a different one favors clearly steepening inner profiles for larger $n$. Indeed we note that the numerical values for the inner slopes obtained by \cite{knollmann_etal_2008} are, for the larger $n$ investigated, in quite good agreement with the exponent predicted by stable clustering. Our results here are for power law initial conditions, but we could equally use the model to study both initial conditions and an appropriately modified time-dependent damping rate mimicking the $\Lambda$CDM model in one dimension. The exact scale invariance in the strongly non-linear regime would certainly be broken in this case, and the various characteristic scales introduced will be imprinted on the clustering. There would thus not be a single correlation exponent, but a slowly varying one. One would certainly expect, however, the qualitative nature of the clustering to be unchanged, just as in three dimensions clustering appears to be qualitatively the same in ``scale-free" cold dark matter cosmologies and $\Lambda$CDM. Further, our considerations here are strictly relevant only to dissipationless {\it cold} dark matter simulations. If the initial conditions are ``warm" or ``hot", or if other non-gravitational interactions are turned on, the associated physical effects will tend to smooth out the matter distribution up to some scale (and thus destroy the scale invariance up to this scale). Nevertheless, if the conclusion B is correct even for this idealized case, it is likely to have very important observational implications relevant to testing standard cosmological models. Intrinsically clumpy or grainy halos lead, for example, to very different predictions for dark matter annihilation (see, e.g. \cite{goerdt_etal_2007, diemand_etal_nature_2008, afshordi_etal_2010})). Concerning the compatibility of such a distribution of cold dark matter with observations of the distribution of visible matter at sub-galactic scales --- which would be expected to lead to an amplified ``missing satellite" problem --- we note that recent studies such as \cite{wadepuhl+springel_2010} show that solid conclusions in this respect will be reached only on the basis of a much improved understanding of the processes involved in star formation. At larger scales the possible link to the striking power-law behavior which characterizes galaxy correlations over several decades in scale (see, e.g., \cite{peebles_1974, sylos_etal_1998, masjedi_etal_2006}) --- which was the motivation for original work using stable clustering to explain such behaviour from power law initial conditions \citep{peebles_1974} and is naturally interpreted as indicative of underlying scale invariance in the matter distribution (see, e.g. \cite{sylos_etal_1998, gaite_2007, antal_etal_2009}) --- is intriguing, and will be discussed elsewhere. We acknowledge useful discussions with Andrea Gabrielli, Jean-Michel Levy, Francesco Sylos Labini and Tirawut Worrakitpoonpon. We thank the anonymous referee for several very useful suggestions.

1012.1515

1012

1012.3383_arXiv.txt

The adaptive optics system at the 3.6 m AEOS telescope was used to measure the astrometry and differential magnitude in \textit{I}-band of 56 binary stars in 2002. The astrometric measurements will be of use for future orbital determination, and the photometric measurements will be of use in estimating the spectral types of the component stars. Two candidate companions were detected, but neither is likely to be gravitationally bound. Nine systems had not been observed in over 40 years. Eight of these are shown to share common proper motion, while HD 182352 is shown to be a background star. One of the two components of the HD 114378 ($\alpha$ Com) is shown to be a variable star of unknown type. In addition, 86 stars were unresolved and the full-width half maxima of the images are presented.

Speckle interferometry has been the dominant technique for observing visual binary stars for the last several decades. There are several groups actively monitoring binary stars with speckle interferometry. One of the longest programs is that of the U.S. Naval Observatory \citep{mason2010}; others include the PISCO program \citep{prieur2010}, the WIYN speckle program \citep{horch2010} and that of Docobo and collaborators \citep{docobo2010}. The main purpose of these observations is measuring the astrometry of the system. These measurements will eventually enable the computation of an orbit for the binary star, which leads to the determination of the dynamical masses of the component stars \citep{docobo2010}. To compute a high quality orbit, it requires many epochs of observations over at least one period of the orbit. For some orbits, this may take hundreds of years. Speckle interferometry systems are relatively inexpensive to build and straightforward to operate. They are able to quickly observe large number of stars; as many as several hundred observations per night. For all of its successes, the technique does have its limitations. It has a fairly low limiting dynamic range. The exact limit depends on the camera used, but ranges from 3-5 magnitudes \citep{mason1994}. Also, in most cases true images are not created, but instead an autocorrelation, which can lead to quadrant ambiguity where the position of the companion is off by 180$^\circ$ \citep{bagnuolo1992}. It is also difficult to extract the photometry from speckle interferograms, especially with the commonly used intensified detectors \citep{roberts1998}. Adaptive optics (AO) solves many of the problems of speckle interferometry. AO systems sense the phase aberrations in the incoming starlight and correct it in real time. Because the AO system corrects the atmospheric distortion, the images from the science camera have a much higher signal-to-noise ratio and can achieve a dynamic range of 10 magnitudes, though this depends on the separation between the star and the companion \citep{turner2008}. With a coronagraph or image subtraction techniques the achievable dynamic range can be increased dramatically \citep{oppenheimer2009}. The other great advantage of AO is that it produces a true image of the system. This removes the quadrant ambiguity that is common in speckle interferometry, it also allows for the measurement of the photometry of the individual stars. This information can be used to estimate the spectral type of the companion \citep{tenBrummelaar2000}. This can then be compared to the mass (or mass sum) computed from the orbit. If the observations are done in multiple filters, the components can be put on a colour-colour diagram \citep{caballero-nieves2010} for additional understanding of the system. Most adaptive optics observations have been focused on specific projects such as multiplicity surveys (e.g. Turner et al. 2008), but there is a great benefit to using AO for long term monitoring of binary stars. The increased dynamic range of AO allows it to be used in the study of systems that the speckle interferometry has been unable to observe. Many of these systems were discovered decades ago with visual methods (e.g. Burnham 1894) and the published astrometry has large errors. Between 2001 and 2006, the Advanced Electro-Optical System (AEOS) telescope and AO system \citep{roberts2002} were used to observe binary stars in \textit{I}-band in order to collect astrometric data to improve orbit determination and to provide photometric data for spectral class determination. This paper presents the measurements from data collected in 2002. The other observations will be be presented in subsequent papers. Most AO systems have science cameras that observe in the near-IR, only a handful of systems have operated in the visible. These include the Mt. Wilson system \citep{shelton1995} and the systems at the U.S. Air Force telescopes at the Starfire Optical Range \citep{fugate1994, spinhirne1998} and on Maui \citep{roberts2002}. During the last decade, only the U.S. Air Force's AEOS telescope on Maui was available for astronomical observations and is currently unavailable for astronomical observations. As such, photometric measurements from the AEOS telescope are unique and unlikely to be repeated in the near future. Photometric measuremetns in the visible are especially useful when combined with near-IR photometric measurements. The addition of visible measurements to near-IR measurements, decreases the uncertainty in spectral classification of stars.\citep{hinkley2010}. Of course, since the binary systems are dynamic, the astrometric measurements can not be duplicated.

The AO system at the 3.6 m AEOS telescope was used to measure the astrometry and differential magnitude in \textit{I}-band of 56 binary stars in 2002. The astrometric measurements will be of use for future orbital determination, and the photometric measurements will be of use in estimating the spectral types of the component stars. In addition, 86 stars had no resolved companions, and for these full-width half maxima of the images are presented. In addition, the observations allow for the analysis of specific stars. Through analysis of differential magnitude from AO and speckle interferometry, it was determined that the HD 114378 system contains a variable star with a relatively short period. Follow up measurements with CCD photometry are needed to determine the exact type of variable star. Simultaneous measurements with AO can be used to identify which component is the variable. The proper motion of nine stars that had not been observed in decades were analysed. HD 182352 was determined to be a background star, while the rest share common proper motion. Two possible new companions were detected to already known binary stars. A candidate companion to HD 10425 was detected and confirmed with follow up observations several years later. While the system does not appear hierarchal, it is possible it is a physical system with a high inclination. Follow up multi-filter near-IR AO observations are needed to determine if it is a physical system through analysis of proper motion and colour photometry. The candidate companion to HD 118889 has a smaller probability of being physical, but follow up observations are needed to determine if the companion is physical or an artifact.

1012.3383

1012

1012.3456_arXiv.txt

We report the preliminary results of a survey for water vapor in a sample of eight C stars with large mid-IR continuum fluxes: V384 Per, CIT 6, V Hya, Y CVn, IRAS 15194-5115, V~Cyg, S~Cep, and IRC+40540. This survey, performed using the { HIFI} instrument on board the {\it Herschel Space Observatory}, entailed observations of the lowest transitions of both ortho- and para-water: the 556.936~GHz $1_{10}-1_{01}$ and 1113.343~GHz $1_{11}-0_{00}$ transitions, respectively. Water vapor was unequivocally detected in all eight of the target stars. Prior to this survey, IRC+10216 was the only carbon-rich AGB star from which thermal water emissions had been discovered, in that case with the use of the {\it Submillimeter Wave Astronomy Satellite (SWAS)}. Our results indicate that IRC+10216 is not unusual, except insofar as its proximity to Earth leads to a large line flux that was detectable with {\it SWAS}. The water spectral line widths are typically similar to those of CO rotational lines, arguing against the vaporization of a Kuiper belt analog (Ford \& Neufeld 2001) being the {\it general} explanation for water vapor in carbon-rich AGB stars. There is no apparent correlation between the ratio of the integrated water line fluxes to the 6.3 micron continuum flux -- a ratio which measures the water outflow rate -- and the total mass-loss rate for the stars in our sample.

Over the past two decades, observations of the envelopes of AGB (asymptotic giant branch) stars have revealed a molecular composition that is still not entirely understood. In early models for the chemistry within the circumstellar envelopes of AGB stars, the inner envelope was assumed to possess a composition determined by thermochemical equilibrium (TE) within the stellar photosphere, while the outer envelope, exposed to the interstellar radiation field, exhibits a time-dependent gas-phase chemistry that is driven by the photodissociative effects of ultraviolet radiation (e.g.\ Glassgold 1996). In TE, the carbon-to-oxygen ratio is crucial in determining the photospheric composition; cool oxygen-rich stars, with $\rm C/O < 1$, are expected to have photospheres that are dominated by CO and H$_2$O, while those of carbon-rich stars will be dominated by CO and C$_2$H$_2$. However, a variety of observations have revealed anomalous abundances for several molecules that cannot be explained by models in which the inner envelope entirely reflects the photospheric abundances in TE and the ensuing chemistry proceeds in the gas phase. More recent models involving shocks (Willacy \& Cherchneff 1998; Cherchneff 2006), grain surface reactions (e.g. Willacy 2004), the vaporization of orbiting objects (Ford \& Neufeld 2001), or the penetration of UV radiation through a clumpy outflow to the inner envelope (Decin et al.\ 2010) have been variously invoked to explain the observed abundances of such species as CO, HCN, SiO, CS, H$_2$CO, OH and H$_2$O in the envelopes of both O-rich and C-rich AGB stars. \begin{deluxetable*}{lcrccrcccccc} \tablewidth{0pt} \tabletypesize{\scriptsize} \tablecaption{Sample of C stars surveyed for water vapor} \tablehead{Star & R.A. & Dec. & D & $\dot M$ & $v_{\rm sys}$ & \multicolumn{2}{c}{Date (2010)} & \multicolumn{2}{c}{Duration (s)} & \multicolumn{2}{c}{Noise (mK) $^f$}\\ & (J2000) & (J2000) & (pc) & (a) & (b) & 557 & 1113 & 557 & 1113 & 557 & 1113 \\ &&&&&& GHz & GHz & GHz & GHz & GHz & GHz } \startdata V384 Per & 03 26 29.5 & +47 31 50 & 740 & $5.5$ & --16.2 & Sep 1 & Sep 2 & 3972$^d$ & 987 & 2.4 & 14\\ % IRC+10216$^c$ & 09 47 57.4 & +13 16 44 & 110 & $15$ & --25.5 & May 4,11 & May 11 & 3150$^d$ & 3200$^a$ & 2.7 & 14 \\ % CIT 6$^c$ & 10 16 02.2 & +30 34 19 & 320 & $1.9$ & --1.6 & May 11 & May 11 & 4719\phantom{$^d$} & 987 & 1.9 & 16 \\ % V Hya & 10 51 37.2 & --21 15 00 & 600 & $ 2.1$ & --17.0 & Jun 23 & Jun 9 & 3972$^d$ & 987 & 3.4 & 15 \\ % Y CVn & 12 45 07.8 & +45 26 25 & 220 & $0.15$ & 21.1 & Jun 23 & Jun 9 & 3972$^d$ & 987 & 2.2 & 16 \\ % IRAS 15194-5115 & 15 23 04.9 & --51 25 59 & 590 & $18$ & --15.0 & Aug 2 & N/A$^e$ & 3972$^d$ & N/A$^e$ & 2.1 & N/A$^e$\\ V Cyg & 20 41 18.3 & +48 08 29 & 530 & $2.5$ & 14.3 & Jun 23 & Mar 5 & 3972$^d$ & 987 & 2.6 & 16 \\ % S Cep & 21 35 12.8 & +78 37 28 & 720 & $3.4$ & --15.5 & N/A$^e$ & May 12 & N/A$^e$ & 987 & N/A$^e$ & 13 \\ % IRC+40540$^c$ & 23 34 27.7 & +43 33 02 & 780 & $7.0$ & --17.0 & Jun 23 & Jun 9 & 3972$^d$ & 987 & 2.5 & 15\\ % \enddata \tablenotetext{a}{Mass loss rate in units of $10^{-6}\,M_\odot\,\rm yr^{-1}$} \tablenotetext{b}{Systemic velocity in km s$^{-1}$ relative to the local standard of rest} \tablenotetext{c}{IRC+10216, CIT6, and IRC+40540 are known also as CW Leo, RW LMi, and LP And respectively.} \tablenotetext{d}{Divided into two observations of equal duration with slightly different LO settings (see text)} \tablenotetext{e}{N/A = not available; planned observation not yet carried out} \tablenotetext{f}{R.m.s.\ in a 1 km s$^{-1}$ bandwidth} \end{deluxetable*} While most of the principal circumstellar molecules are easily detected by means of ground-based observations, water has been more elusive. Although water maser emissions from oxygen-rich stars can be detected from the ground, thermal water emissions are detectable only with satellite observatories. As expected, the {\it Infrared Space Observatory (ISO)} detected luminous water emissions from a variety of oxygen-rich AGB stars (e.g. Barlow 1999, Neufeld et al.\ 1999). More surprisingly, later observations performed (Melnick et al.\ 2001) with the {\it Submillimeter Wave Astronomy Satellite (SWAS)} -- and subsequently confirmed with the ODIN satellite (Hasegawa et al.\ 2006) -- have led to the discovery of water vapor in the envelope of the nearest carbon-rich AGB star { with a massive envelope}, IRC+10216. Prior to the recent launch of the {\it Herschel Space Observatory} (Pilbratt et al.\ 2010), this detection of the lowest, $1_{10} - 1_{01}$, rotational transition of ortho-water at 556.936 GHz by SWAS and ODIN remained the only case in which thermal water vapor emission had been detected from a C star. The {\it Herschel Space Observatory} presents an opportunity to search for water vapor in C stars with unprecedented sensitivity; with a primary mirror of collecting area $\sim 30$ times that of the SWAS mirror, and with its Heterodyne Instrument for the Far Infrared (HIFI; de Graauw et al. \ 2010) employing cryogenic mixers of much lower noise temperature than those flown on SWAS, {\it Herschel} makes it { feasible} to search for water in AGB stars considerably more distant than IRC+10216. In addition, {\it Herschel} provides access to multiple transitions of water vapor. As part of the HIFISTARS Key Program, we carried out a survey for water vapor in a sample comprising eight C stars in addition to IRC+10216 -- V384 Per, CIT 6, V Hya, Y CVn, IRAS 15194-5115, V Cyg, S Cep, and IRC+40540 -- with the goal of determining whether IRC+10216 is unusual (in addition to being nearby), or whether water vapor is widely detectable in the envelopes of C stars. Our survey involves observations of the lowest transitions of both ortho- and para-water, the 556.936~GHz $1_{10}-1_{01}$ and 1113.343~GHz $1_{11}-0_{00}$ transitions respectively, toward a selection of C stars with large mid-IR continuum fluxes. Table 1 lists the sources in our sample, together with estimates of their distances, $D$, mass-loss rates, ${\dot M}$, and systemic velocities relative to the local standard of rest. { Although the quantity ${\dot M}/D^2$ is fairly well constrained, uncertainties in $D$ make both parameters quite uncertain} -- a fact that is demonstrated by the wide dispersion of estimates given in the literature -- but in the interest of uniformity we adopt the values tabulated by Groenewegen et al.\ (2002), supplemented by those of Sch\"oier \& Olafsson (2001) for the case of Y CVn (for which estimates were not given by Groenewegen et al.\ 2002). With the exception of Y CVn, which is classified as a J-type variable carbon star of spectral type CV5J (with V indicating variability and J indicating a large $\rm ^{13}C/^{12}C$ ratio), all the stars in the sample are of spectral type CV6 or CV7. V Hya is also of a somewhat different character than the other sources; { it is a semi-regular variable (SRa), quite probably a binary (Sahai et al.\ 2003), and possesses} a high velocity outflow and bipolar symmetry that suggests it is in an early stage of {\it post}-AGB evolution (Kahane et al.\ 1996, Knapp et al.\ 1997). { The mass-loss rate given in Table 1 is for the component with outflow speed $\sim 15\,\rm km\,s^{-1}$ (identified as the ``normal slow wind'' by Knapp et al.\ 1997).} In this {\it Letter}, we report preliminary results of this survey, based on the observations performed to date toward the eight target stars. The detection of the $1_{11}-0_{00}$ water transition toward V Cygni in the very first observation performed in this survey has been reported previously (Neufeld et al.\ 2010). The data acquisition and reduction are described in \S 2, and the results are presented in \S 3. In \S 4, the results are discussed with reference to various scenarios that have been proposed for the origin for water vapor around C stars. \begin{figure} \includegraphics[scale=0.9]{fig1.eps} \caption{Black histogram: Herschel/HIFI spectra of the $1_{10}-1_{01}$ 556.936~GHz transition of ortho-H$_2$O (left) and the $1_{11}-0_{00}$ 1113.343~GHz transition of para-H$_2$O (right) observed toward V384 Per, IRC+10216, CIT6, and V Hya. Blue histogram: analogous spectra for CO $J= 10-9$, where available, scaled to the same peak antenna temperature. The spectra have been continuum-subtracted and rebinned to a spectral resolution of 1~km~s$^{-1}$. Doppler velocities are expressed in km~s$^{-1}$ relative to the systemic velocity of the source, and antenna temperatures (vertical axis) are shown in mK. Dotted lines indicate expansion velocities from the literature ({ Groenewegen et al.\ 2002).}} \end{figure} \begin{figure} \includegraphics[scale=0.9]{fig2.eps} \caption{Same as Fig. 1, for Y CVn, IRAS-15194, V Cyg, S Cep and IRC+40540. For Y CVn and IRAS-15194 the blue histogram shows CO $J=6-5$ instead of $J=10-9$} \end{figure}

Although there are real differences between the H$_2$O and CO line profiles for some of the spectra shown in Figures 1 and 2 (most notably for the IRC+10216 spectra which have the highest signal-to-noise ratio), the line widths are typically quite similar. In particular, the half width at zero intensity is typically close to literature estimates of the expansion velocity (outer dotted vertical lines in each panel). This feature of the spectra indicates that water-emitting gas is typically present over a wide range of outflow directions, including streamlines directed close to the line-of-sight, and argues strongly against the vaporization of a Kuiper belt analog (Ford \& Neufeld 2001) being the {\it general} explanation for water vapor in carbon-rich AGB stars. Any collection of orbiting icy objects with a flattened structure (like the Kuiper belt) would release water that would be entrained in the outflow and carried along streamlines with small inclinations to the equatorial plane of the structure. Viewed nearly edge-on, such an anisotropic distribution of water vapor could indeed yield a spectral line of width comparable to that of the CO lines; but viewed nearly face-on, the water lines would be considerably narrower. Thus, while the vaporization of a Kuiper belt might account for the water line profile observed in any {\it individual} source, the statistical ensemble now available from our survey is not consistent with Kuiper belt vaporization as a general explanation for water vapor. (Note, however, that based upon multitransition observations available for IRC+10216 but not for the other sources in our sample, Decin et al.\ 2010, and subsequently Neufeld et al.\ 2011, have argued that the distribution of water vapor in that source is inconsistent with the vaporization of small icy objects on circular orbits being the source of the water in that source.) \begin{figure} \includegraphics[scale=0.50]{fig3.eps} \caption{Ratio of integrated line flux to 6.3$\mu$m continuum flux for the 557 GHz (black diamonds) and 1113 GHz (red asterisks) water lines, as a function of the estimated total mass-loss rate (Table 1)} \end{figure} Detailed calculations by GNM predict luminosities for the $1_{10}-1_{01}$ and $1_{11}-0_{00}$ transitions that are roughly proportional to the 6.3 $\mu$m continuum luminosity, and that the ratio of integrated line flux to 6.3$\mu$m continuum flux is an increasing function of the water mass-loss rate. Because collisional excitation is negligible except in the densest CSEs, the line fluxes show almost no dependence upon the gas density; thus the ratio of integrated line flux, $F_l$, to continuum flux, $F_{6.3}$, is nearly independent of the total mass-loss rate. In Figure 3, we plot the $F_l/F_{6.3}$ ratio for each star as a function of the total mass-loss rate, ${\dot M}$. Here, we used the results tabulated in Table 2, with integrated line fluxes in units of $10^{-17}$ W m$^{-2}$, and the 6.3 $\mu$m continuum fluxes are in units of $10^3$ Jy; thus the plotted ratio has units of MHz. The plotted values of ${\dot M}$ are the literature estimates tabulated in Table 1, { with error bars indicating an estimated uncertainty of 0.2 dex}. If the water vapor present in these CSEs is produced by chemical means within the outflow with some universal H$_2$O/H$_2$ ratio, the $F_l/F_{6.3}$ ratio (measuring water outflow rate) and the total mass-loss rate would show a positive correlation. There is no discernable correlation in Figure 3, suggesting that stars with larger mass outflow rates tend to have smaller water {\it abundances} (although it remains unclear whether this tendency indicates a {\it causative} relationship). Quantitative estimates of the water outflow rates and abundances will require detailed modeling that will be the subject of a future paper (and that we hope to constrain further with future {\it Herschel} observations.) For the two sources that have been studied in detail to date, V Cyg (Neufeld et al.\ 2010) and IRC+10216 (GNM; Neufeld et al.\ 2011), the derived water outflow rates support the lack of positive correlation suggested by Figure 3. While IRC+10216 has a total mass-loss rate $\sim 10$ times that of V Cyg, it shows a water outflow rate that is only $\sim$ one-third as large; { the inferred H$_2$O/H$_2$ abundance ratio in the CSE of V Cyg ($2 \times 10^{-6}$) is therefore $\sim 25$ times larger than that inferred for IRC+10216 ($8 \times 10^{-8}$).}

1012.3456

1012

1012.0455_arXiv.txt

Gamma-ray emission produced by interactions between cosmic rays (CRs) and interstellar gas traces the product of their densities throughout the Milky Way. The outer Galaxy is a privileged target of investigation to separate interstellar structures seen along the line of sight. Recent observations by the \textit{Fermi} Large Area Telescope (LAT) shed light on open questions of the EGRET era about the distribution of CR densities and the census of the interstellar medium. The gradient of \g-ray emissivities measured in the outer Galaxy is significantly flatter than predictions from widely used CR propagation models given the rapid decline of putative CR sources beyond the solar circle. Large propagation volumes, with halo heights up to 20~kpc, or a flat CR source distribution are required to match the data. Other viable possibilities include non-uniform CR diffusion properties or more gas than accounted for by the radio/mm-wave data. \g-ray data constrain the evolution of the $\xco=\nhd/\wco$ ratio within a few kpc from the Sun. There is a significant increase by a factor~2 from nearby clouds in the Gould Belt to the local spur. No further significant variations are measured from the local spur to the Perseus spiral arm. At the level of statistical accuracy provided by the LAT data, the most important source of uncertainty, often overlooked so far, is due to the optical depth correction applied to derive the column densities of $\hi$. Reliable determinations of the amount of atomic gas in the plane are key to better probe the properties of CRs in the Galaxy.

Interstellar \g-ray emission is produced by interactions of high-energy cosmic rays (CRs) with the gas in the interstellar medium (ISM) and the soft interstellar radiation fields. Its observations carry information about CR properties in distant locations and it provides a tracer of the total interstellar gas densities to be compared with radio/mm-wave data: the 21-cm line of atomic hydrogen, $\hi$, and the 2.6-mm line of CO, used as a surrogate tracer of molecular mass. Open issues in the understanding of the interstellar \g-ray emission concern the identification and spatial distribution of CR sources and the census of the ISM, notably the $\xco=\nhd/\wco$ conversion factor. The outer Galaxy is a privileged observational target since the Doppler shift of radio/mm-wave lines due to the Galactic rotation unambiguously locates the emitting clouds. There are two longitude windows with a steep velocity gradient leading to a good kinematic separation \cite{ref:2}. We reported analyses of recent measurements by the Large Area Telescope (LAT) on board the \textit{Fermi \g-ray Space Telescope} \cite{ref:4} for the second \cite{ref:3} and third \cite{ref:5} Galactic quadrants. The component separation based on likelihood fitting allowed us to extract the emissivities per H atom, $\qhi$, and per $\wco$ unit, $\qco$, in several regions along the lines of sight as described in table~\ref{tab:regions}. We refer the interested reader to the aforementioned papers for details about the analysis and we briefly discuss here the implications of the results for the distribution of CRs in the Galaxy and the calibration of the $\xco$ ratio. \begin{table}[!htbp] \caption{Regions seen toward the outer Galaxy and their approximate Galactocentric distances.} \label{tab:regions} \begin{tabular}{rcc} \hline & second quadrant & third quadrant \\ & $100^{\circ}<l<145^{\circ}$ & $210^{\circ}<l<250^{\circ}$ \\ \hline Gould Belt & $8.5-8.8$ kpc & -- \\ Local Arm & $8.8-10$ kpc & $8.5-10$ kpc \\ interarm region & -- & $10-12.5$ kpc \\ Perseus arm & $10-14$ kpc & $12.5-16$ \\ outer region & $>14$ kpc & $>16$ kpc \\ \hline \end{tabular} \end{table}

1012.0455

1012

1012.2450_arXiv.txt

{It is not known how many globular clusters may have been left undetected towards the Galactic bulge.} {One of the aims of the VISTA Variables in the Via Lactea (VVV) Survey is to accurately measure the physical parameters of the known globular clusters in the inner regions of the Milky Way and to search for new ones, hidden in regions of large extinction.} {Deep near infrared images give deep $JHK_{\rm S}$-band photometry of a region surrounding the known globular cluster UKS~1 and reveal a new low-mass globular cluster candidate that we name VVV CL001. } {We use the horizontal branch red clump in order to measure E(B-V)$\sim$2.2 mag, $(m-M)_0=16.01$ mag, and D=15.9 kpc for the globular cluster UKS~1. Based on the near-infrared colour magnitude diagrams, we also measure that VVV CL001 has E(B-V)$\sim$2.0, and that it is at least as metal-poor as UKS~1, however, its distance remains uncertain.} {Our finding confirms the previous projection that the central region of the Milky Way harbors more globular clusters. VVV~CL001 and UKS~1 are good candidates for a physical cluster binary, but follow-up observations are needed to decide if they are located at the same distance and have similar radial velocities. }

\label{Intro} The inner regions of the Milky Way have been mapped thoroughly at all wavelengths. Yet, it is not known if there are still some distant globular clusters awaiting to be discovered, hidden beyond the bulge, due to the high density of stellar sources and the large and inhomogeneous interstellar extinction. Near-IR surveys have an advantage for searching these regions. Indeed, the 2MASS discovered two new globular clusters \citep{Hurt}. But the limiting magnitude of 2MASS \citep[$K_{\rm S}$=14.3, for 10$\sigma$-detections;][]{Skrutskie06} may prevent the discovery of fainter objects, especially if they are located in highly reddened regions. The asymmetry of the spatial distribution of known globular clusters around the Galactic center indicates that previous observations may have overlooked some additional globular clusters. \cite{Ivanov2005} recently estimated that there may be about 10 clusters missing towards the inner Milky Way. The recent discoveries (last 10 years) include both faint (low mass) halo clusters as well as reddened globular clusters projected toward the bulge, e.g. 2MASS~GC01 and 2MASS~GC02 by \cite{Hurt} \citep[see also][]{Ivanov2000}, ESO 280~SC06 by \cite{Ortolani2000}, GLIMPSE~C01 by Kobulnicky et al. (2005; but see also Ivanov et al. 2005; Davies et al. 2010), GLIMPSE~C02 by \cite{Kurtev}, AL-3 by \cite{Ortolani2006}, FSR 1735 by \cite{Froebrich}, Koposov~1 and Koposov~2 by \cite{koposov07}, FSR~1767 by \cite{Bonatto}, Whiting~1 \citep[][]{carraro05} and Pfleiderer~2 by \cite{Ortolani2009}. The VISTA Variables in the Via Lactea (VVV) Public Survey has started mapping the inner disk and bulge of our Galaxy with VISTA 4m telescope (Visible and Infrared Survey Telescope for Astronomy) in the near-IR \citep{Minniti,Saito}. One of the main scientific goals of the VVV Survey is to study the bulge globular clusters and to search for new clusters. Here we present VVV~CL001, the first globular cluster candidate discovered by the VVV Survey.

The VVV Survey is searching for missing globular clusters in the inner regions of the Milky Way galaxy. We report here the discovery of VVV CL001 (Figure \ref{fig1}), a low-mass globular cluster candidate at $\alpha=17:54:42.5$, $\delta=-24:00:53$. This is located only 8 arcmin away from the Galactic globular cluster UKS~1 in the sky. We take advantage of this spatial proximity to use UKS~1 as reference in order to estimate the parameters of this newly discovered cluster. We measure the distance and reddening of UKS~1, finding $E(B-V)_{UKS1}=2.5$ and $D_{UKS1}=15.9$kpc. We present the first colour-magnitude diagrams of the new cluster VVV~CL001 (Figure \ref{cmd}), estimating $E(B-V)= 2.0$. We cannot define the mean $K_{\rm S}$-band magnitude of the horizontal branch red clump for this cluster, and therefore its distance is uncertain. A very crude estimate by comparison with the RGB of UKS~1 gives a similar distance, placing VVV CL001 well beyond the bulge of the Milky Way. Observations in the $K_{\rm S}$-band to be acquired in the following seasons by the VVV Survey would allow us to improve this CMD of Figure \ref{cmd} and define the HB as well as the MSTO, which would allow a photometric age determination. Also, we estimate that the RR Lyrae of VVV~CL001 would be within the limit of detection of our VVV Survey. We cannot definitely conclude that the proximity of UKS~1 and VVV CL001 on the sky imply that they compose a binary cluster because the distance to VVV CL001 is too uncertain. This remains as an exciting possibility that needs to be confirmed not only by means of a more accurate distance determination, but also by measuring their respective radial velocities from spectroscopic measurements. Finally, the present results are very encouraging and we conclude that the VVV Survey can potentially provide the largest and most homogeneous census of globular clusters in the survey area, out to well beyond the Galactic center, even in regions of large extinction. \begin{figure*}[!ht] \centering \includegraphics[width=7.5cm]{Minniti_fig1a.jpg} \includegraphics[width=4.7cm]{Minniti_fig1b.jpg} \caption{Discovery image (left panel, $JHK_{\rm S}$ colour composite) showing VVV~CL001 on the left and UKS~1 on the right. The cluster UKS~1 was discovered by \cite{Malkan80}, with a total magnitude $I<16$ mag, and it was invisible in Palomar Survey plates, (and up to a few years ago UKS~1 was known as "the faintest globular cluster of the Milky Way"). The right panel shows a zoom into the region centered VVV~CL001, corresponding to the white box (1\farcm75 x 1\farcm5).} \label{fig1} \end{figure*} \begin{figure}[!ht] \includegraphics[width=8.5cm]{Minniti_fig2.jpg} \caption{J-band image of the VVV~CL001 region (left) and density of all objects. The contours in the left panel refer to the intensity level of the image. In contrast, to derive the significance of the over-density (right) only detections with a Dophot classification parameter 1 or 7 \citep{Schechter} were selected. The sources were binned into 60x60 pixel bins (20 by 20 arcsec). We note that evenwith PSF photometry we are not able to resolve the inner region of VVV~CL001 (or UKS~1), which certainly affects the numerical results of our analysis (density, significance of the detected over-density). Both plots show the same 272 by 272 arcsec section of the J-band image. } \label{density} \end{figure} \begin{figure}[!ht] \includegraphics[height=4.5cm]{Minniti_fig3.jpg} \caption{Discovery near-infrared colour-magnitude diagram of VVV~CL001 for a 15 arcsec radius area centered on the cluster (left) compared with a field of same area located $3\farcm75$ east (middle), and with a field of UKS~1 of the same size and offset the cluster center by (45\farcs0) (right). The small offset for UKS~1 was applied, since the center of this cluster is barely resolved. The horizontal line at $Ks=14.3$ shows the limit of 2MASS photometry. Comparing UKS~1 and VVV~CL001 we find that the RGB of UKS~1 is much more populated, and hence VVV~CL001 to be far less massive then UKS~1.} \label{cmd} \end{figure} \begin{figure}[!ht] \includegraphics[height=4.5cm]{Minniti_fig4.jpg} \caption{CMD of VVV~CL001 (diamonds) compared with UKS~1 (dots). The horizontal branch red clump of UKS~1 is seen at $K_{\rm S}=15.3$, $J-K_{\rm S}=2.0$. This diagram shows that the RGB of VVV~CL001 is slightly bluer than that of UKS~1, from which one can conclude that VVV~CL001 is less reddened or more metal-poor than UKS~1, and that the RGB of UKS~1 is much more populated, indicating the low mass of VVV~CL001. The legend gives the radius of the selected cluster region, whereas the arrows indicate the reddening, derived for UKS~1.} \label{fig4} \end{figure} \begin{table} \caption[]{Properties of VVV~CL001, based on the VVV data, compared to two known extremely low luminosity globular clusters Koposov~1 \& ~2 \citep[][]{koposov07}.\\ $^{*}$based on a distance of 16 kpc and a selected radius of 15 arcsec\\ } \label{cluster} \vskip 0.2cm \begin{tabular}{l r r r} \hline \noalign{\smallskip} & VVV CL001 & Koposov~1 & Koposov~2 \\ Position [l,b] & 4.99$^\circ$, 1.84$^\circ$ & 260.98$^\circ$, 70.75$^\circ$ & 195.11$^\circ$, 25.55$^\circ$ \\ Distance [kpc] & uncertain & $\approx$ 50kpc & $\approx$ 40 kpc\\ Radius $[arcmin]$ & $\approx$1~pc $^{*}$& $\approx$~3~pc & $\approx$~3~pc\\ \noalign{\smallskip} \hline \end{tabular} \end{table}

1012.2450

1012

1012.2385_arXiv.txt

{ Monte Carlo techniques have been used to evaluate the statistical and systematic uncertainties in the helium abundances derived from extragalactic H~II regions. The helium abundance is sensitive to several physical parameters associated with the H~II region. In this work, we introduce Markov Chain Monte Carlo (MCMC) methods to efficiently explore the parameter space and determine the helium abundance, the physical parameters, and the uncertainties derived from observations of metal poor nebulae. Experiments with synthetic data show that the MCMC method is superior to previous implementations (based on flux perturbation) in that it is not affected by biases due to non-physical parameter space. The MCMC analysis allows a detailed exploration of degeneracies, and, in particular, a false minimum that occurs at large values of optical depth in the He~I emission lines. We demonstrate that introducing the electron temperature derived from the [O~III] emission lines as a prior, in a very conservative manner, produces negligible bias and effectively eliminates the false minima occurring at large optical depth. We perform a frequentist analysis on data from several ``high quality'' systems. Likelihood plots illustrate degeneracies, asymmetries, and limits of the determination. In agreement with previous work, we find relatively large systematic errors, limiting the precision of the primordial helium abundance for currently available spectra. }

Standard big bang nucleosynthesis (SBBN) using the baryon density determined by WMAP \citep{wmap,wmap10} predicts the initial abundances of D, $^{3}$He, $^{4}$He, and $^{7}$Li \citep{cfo,coc,coc2,cyburt,coc3,cuoco,serp,cfo5}, allowing one to probe the early universe at redshifts of order $10^{10}$ \citep{wssok,osw,fs}. Therefore, the observed abundances provide a valuable check on the theory of SBBN, its concordance with the measurements of the microwave background radiation, and the content and interactions of the universe during the period of BBN \citep{MM93,sar,cfos}. To test these predictions, the observed abundances must be determined with relatively high precision. Because of the logarithmic relationship between the baryon to photon ratio, $\eta$, and the primordial helium abundance, Y$_{p}$, the uncertainty of Y$_{p}$ must be $<1\%$ to meaningfully test the theory. The 7-year WMAP value for $\eta$ is $(6.19 \pm 0.15) \times 10^{-10}$, \citet{wmap10}. For comparison, the SBBN calculation of \citet{cfo5}, assuming the WMAP $\eta$ and a neutron mean life of $885.7 \pm 0.8$ s \citep{rpp}, yields $Y_p = 0.2487 \pm 0.0002$, a relative uncertainty of only 0.08\%. The determination of Y$_{p}$ is facilitated through low metallicity H~II regions in dwarf galaxies. By fitting the helium abundance versus metallicity, one can extrapolate back to very low metallicity, corresponding to the primordial helium abundance \citep{ptp74}. The oxygen to hydrogen ratio, O/H, commonly serves as a proxy for metallicity. Though this area of research has benefited from three decades of development, the determinations of Y$_{p}$ have suffered from significant differences between the results. The difficulties in calculating an accurate and precise measure of the primordial helium abundance are well established \citep{os01,os04,its07}. Here, we introduce a new method based on Markov Chain Monte Carlo (MCMC) techniques. Observations of the helium to hydrogen emission line ratio from extragalactic H~II regions provide a measure of the helium to hydrogen ratio, y$^{+}={n(He~II) \over n(H~II)}$. Correspondingly, the statistical measurement errors in the helium and hydrogen emission line fluxes contribute to the uncertainty on y$^{+}$. Unfortunately, this calculation of y$^{+}$ and its uncertainty are complicated by a myriad of systematic effects. Interstellar reddening, underlying stellar absorption, radiative transfer, and collisional corrections alter the observed flux, complicating the measurement of y$^{+}$, and amplifying the uncertainty. The photons are scattered by dust on their journey (interstellar reddening). The stellar continuum juxtaposes absorption features under nebular emission lines (helium and hydrogen underlying absorption). The H~II region itself absorbs and re-emits photons (radiative transfer); both recombination and collisional excitation contribute to the emission (collisional corrections for helium and hydrogen). None of these processes can be directly measured and, therefore, cannot be removed independent of the observed emission lines and theoretical models. As a result, the uncertainty on y$^{+}$ must reflect the presence of, and lack of certainty regarding, these systematic effects. Determining y$^{+}$ in conjunction with the robust estimation of the model parameters used in correcting for the listed systematic effects requires ``high quality'' spectra. This desired confidence weighs against the need for larger sample sizes to decrease the uncertainty on Y$_p$ (and $dY/dZ$). The importance of Monte Carlo techniques was demonstrated in a ``self-consistent'' analysis method, stemming from the work of \citet{itl94} and \citet{ppr00}, for determining the nebular helium abundance based upon six helium and three hydrogen lines \cite{os01,os04}. In preceding work \cite{AOS}, hereafter AOS, we updated and extended the physical model and integrated the helium and hydrogen calculations with the goals of improving accuracy and removing assumptions. The focus of the current paper is the exploration of a new technique based on a Markov Chain Monte Carlo (MCMC) analysis and departs from the ``self-consistent'' method. Rather than fitting the parametric inputs to a helium abundance, the frequentist approach developed here builds a global likelihood function for all parameters including the helium abundance. As we will demonstrate, the MCMC method is statistically superior to previous efforts, it is more direct and transparent, and it maintains efficiency. Of primary benefit, the results are more rigorous: the solution remains unbiased by the procedure, the uncertainty captures the confidence of the model and measurements, visualization of the parameter space topology reveals the reliability of the determination, and spectra failing to resolve their physical environments are identified. Section \ref{ParameterSimulation} discusses the determination of parameter uncertainties and details the differences in our approach between AOS and this work. The utility of Markov Chains and the computational implementation are described in \S \ref{MCMC}. In \S \ref{Synthetic}, we describe tests using synthetic data to illustrate the method and its utility, with particular emphasis on secondary minima and the incorporation of a temperature prior. MCMC is implemented in analyzing the dataset used in AOS in \S \ref{Galaxies}, and, in \S \ref{Results}, Y$_p$ is determined. Finally, \S \ref{Conclusion} offers a discussion of the exploration and results as well as the next steps in better determining the primordial helium abundance.

\label{Conclusion} The primary result of this work is the demonstration of a statistically rigorous method for determining the uncertainty of the abundance and model parameters. The use of MCMC allows one to efficiently sample the parameter space so that the uncertainties can be calculated directly from the change in $\chi^{2}$ as the parameters are varied from the best-fit solution. Computationally, MCMC is efficient and straightforward. Beyond the improvement in the approach itself, the constructed likelihood distributions for the parameters are instructive in evaluating the quality of the object and illuminating the sources of differences with previous analyses. Particularly illustrative of the benefits of the likelihood approach was the discovery of the increased degeneracy at large optical depth. With synthetic data, a second minimum emerges as the optical depth increases. As this false minimum becomes more significant, the reliability and quality of the object is undermined. In concordance with this predicted effect, several galaxies, each with large optical depth, exhibit prominent second minimums. However, a second benefit of the approach is the straightforward incorporation and interpretation of priors. This allows the electron temperature defined by the [O~III] emission lines to be utilized to eliminate low temperature, unphysical secondary minimums. Therefore, after taking the [O~III] measurement into account, large optical depth objects are well behaved. It also worthy of note that the prior is used very conservatively. This ensures that the solved value of the temperature primarily reflects the helium defined temperature, thus protecting the abundance from any significant bias. Thus, the new MCMC method is a distinct improvement, resulting in a statistically more accurate determination of the helium abundance, the physical parameters associated with the HII region, and their uncertainties. Nevertheless, we found relatively large uncertainties in the helium abundance determinations of individual low metallicity HII regions. This, however, is an indication of the true uncertainty in the measurement and the challenge posed. Future work will investigate the possibilities for improving the result through the use of a revised and expanded set of objects.

1012.2385

1012

1012.0039_arXiv.txt

The densities in the outer regions of clusters of galaxies are very low, and the collisional timescales are very long. As a result, heavy elements will be under-ionized after they have passed through the accretion shock. We have studied systematically the effects of non-equilibrium ionization for relaxed clusters in the $\Lambda$CDM cosmology using one-dimensional hydrodynamic simulations. We found that non-equilibrium ionization effects do not depend on cluster mass but depend strongly on redshift which can be understood by self-similar scaling arguments. The effects are stronger for clusters at lower redshifts. We present X-ray signatures such as surface brightness profiles and emission lines in detail for a massive cluster at low redshift. In general, soft emission (0.3--1.0~keV) is enhanced significantly by under-ionization, and the enhancement can be nearly an order of magnitude near the shock radius. The most prominent non-equilibrium ionization signature we found is the \ion{O}{7} and \ion{O}{8} line ratio. The ratios for non-equilibrium ionization and collisional ionization equilibrium models are different by more than an order of magnitude at radii beyond half of the shock radius. These non-equilibrium ionization signatures are equally strong for models with different non-adiabatic shock electron heating efficiencies. We have also calculated the detectability of the \ion{O}{7} and \ion{O}{8} lines with the future {\it International X-ray Observatory} ({\it IXO}). Depending on the line ratio measured, we conclude that an exposure of $\sim$130--380~ksec on a moderate-redshift, massive regular cluster with the X-ray Microcalorimeter Spectrometer (XMS) on the {\it IXO} will be sufficient to provide a strong test for the non-equilibrium ionization model.

\label{ion_sec:intro} Clusters of galaxies are very sensitive probes to cosmological parameters \citep[e.g.,][]{ARS08,Vik+09}. Systematic uncertainties in precision cosmology using galaxy clusters can be minimized by restricting the sample of clusters to the highest degree of dynamical relaxation, and hence studying relaxed clusters is particularly important. In addition, the outer envelopes of clusters have been thought to be less subjected to complicated physics such as active galactic nucleus feedback, and hence these outer regions may provide better cosmological probes. Clusters are believed to be formed by the continual merging and accretion of material from the surrounding large-scale structure. Based on high resolution $N$-body simulations, it has been found that cluster growth can be divided into an early fast accretion phase dominated by major mergers, and a late slow phase dominated by smooth accretion of background materials and many minor mergers \citep{WBP+02,ZJM+09}. Accretion shocks (or virial shocks) are unambiguous predictions of cosmological hydrodynamic simulations. For the most relaxed clusters with roughly spherical morphology in the outer regions, these simulations show that large amounts of material are accreted through filamentary structures in some particular directions, and more spherically symmetrically in other directions. Most of the simulations assume the intracluster medium (ICM) is in collisional equilibrium. However, because of the very low density in the cluster outer regions ($\ga R_{200}$\footnote{$R_{\Delta}$ is the radius within which the mean total mass density of the cluster is $\Delta$ times the critical density of the universe. The virial radius $R_{\rm vir}$ is defined as a radius within which the cluster is virialized. For the Einstein-de Sitter universe, $ R_{\rm vir} \approx R_{178}$, while for the standard $\Lambda$CDM Universe, $ R_{\rm vir} \approx R_{95}$.}), the Coulomb collisional and collisional ionization timescales are comparable to the age of the cluster. It has been pointed out that electrons and ions there may be in non-equipartition \citep{FL97,Tak99,Tak00,RN09,WS09}, and also the ionization state may not be in collisional ionization equilibrium \citep{YYS+03,CF06,YS06}. If these non-equilibrium processes are not properly taken into account, the measured properties may be biased in these regions. Studying the outer regions of relaxed clusters not only is valuable in understanding the accretion physics and to test the assumptions concerning plasma physics near the shock regions, but it is also very important to test structure formation theory and to constrain the systematic uncertainties in clusters to be used as precision cosmological probes. Before the launch of the {\it Suzaku} X-ray observatory, physical properties such as temperature have never been constrained with confidence beyond roughly one-half of the shock radius. Most of our understanding of these regions is still based on theoretical models and hydrodynamic simulations. Recently, observations by {\it Suzaku} have constrained temperatures up to about half of the shock radius for a few clusters for the first time to better than a factor of $\sim 2$ \citep{GFS+09,Rei+09, Bas+10,Hos+10}. While the uncertainties are still large, there is evidence that the electron pressure in cluster outer regions may be lower than that predicted by numerical simulations assuming collisional equilibrium. Observations of secondary cosmic microwave background anisotropies with the South Pole Telescope (SPT) and the {\it Wilkinson Microwave Anisotropy Probe} ({\it WMAP}) 7 year data also support these results \citep{Kom+10,Lue+10}. These observational signatures are consistent with electrons and ions in non-equipartition, although it is also possible that the hydrodynamic simulations may simply overestimate the gas pressure. Another possibility is that heat conduction outside the cluster may be reducing the gas pressure \citep{Loe02}. Recently, \citet{WSW10} have shown that cosmological parameters will be biased if non-equilibrium effects (in particular, non-equipartition) are not properly taken into account. More observations of the outer regions of clusters are being done or analyzed. In the future, the proposed {\it International X-ray Observatory} ($IXO$)\footnote{http://ixo.gsfc.nasa.gov/} will have the sensitivity to constrain cluster properties out to the shock radius. Thus, a detailed study of the physics in the outer regions of clusters would be useful. In particular, {\it IXO} will have sufficient spectral resolution to resolve many important X-ray lines, and hence the ionization state of the plasma can be determined. In \citet{WS09}, we have studied in detail the X-ray signature of non-equipartition effects in cluster accretion shock regions. In this paper, we extend our study to include non-equilibrium ionization in our calculations. Non-equilibrium ionization calculations have been considered in a number of cosmological simulations to study the very low density warm-hot intergalactic medium (WHIM) surrounding galaxy clusters \citep{YYS+03,CF06,YS06}. At galaxy cluster scales, similar non-equilibrium ionization calculations have focused on merging clusters \citep{AY08,AY10}. These studies generally agree that in the low density ICM and the WHIM, there are significant deviations from ionization equilibrium, and that the effects on the X-ray emission lines are strong. In this work, we focus on the X-ray signatures of non-equilibrium ionization in the accretion shock regions of relaxed clusters. We also discuss the detectability of non-equilibrium ionization effects with the future {\it IXO}. The paper is organized as follows. Section~\ref{ion_sec:method} describes in detail the physical models and techniques to calculate non-equilibrium ion fractions and X-ray observables, which includes the hydrodynamic models we used (Section~\ref{ion_sec:hydro}), and the ionization and spectral calculational techniques (Sections~\ref{ion_sec:IonFracCal}--\ref{ion_sec:XrayCal}). The overall dependence the non-equilibrium ionization effects on the mass and redshift of the cluster is presented in Section~\ref{ion_sec:b_vs_mass_z}. Calculated non-equilibrium ionization signatures are presented in Section~\ref{ion_sec:signatures}, which includes descriptions of particular models we used to present the results (Section~\ref{ion_sec:models}), X-ray spectra (Section~\ref{ion_sec:spectra}), surface brightness profiles (Section~\ref{ion_sec:sb}), and the ratio of intensities of \ion{O}{7} and \ion{O}{8} lines (Section~\ref{ion_sec:line}). We discuss the detectability of \ion{O}{7} and \ion{O}{8} lines and non-equilibrium ionization diagnostics with {\it IXO} in Section~\ref{ion_sec:detect}. Section~\ref{ion_sec:conclusion} gives the discussion and conclusions. Throughout the paper, we assume a Hubble constant of $H_0 = 71.9~h_{71.9}$~km~s$^{-1}$~Mpc$^{-1}$ with $h_{71.9}=1$, a total matter density parameter of $\Omega_{M,0} = 0.258$, a dark energy density parameter of $\Omega_{\Lambda} = 0.742$, and a cluster gas fraction of $f_{\rm gas}=\Omega_b/\Omega_M=0.17$, where $\Omega_b$ is the baryon density parameter for our cluster models in the standard $\Lambda$CDM cosmology\footnote{\scriptsize http://lambda.gsfc.nasa.gov/product/map/dr3/parameters\_summary.cfm}. The clusters have a hydrogen mass fraction $X=76\%$ for the ICM.

\label{ion_sec:conclusion} Studying the physics in the outer regions of clusters is very important to understand how clusters are formed, how the the intracluster gas is heated, as well as to constrain the formation of large-scale structure. Because of the very low density in cluster outer regions, the collisional timescales are very long and comparable to the cluster age. Electrons and ions passed through the accretion shocks may not have enough time to reach equipartition and the ions may be under-ionized. In a previous paper \citep{WS09}, we have studied the non-equipartition effects on clusters using one-dimensional hydrodynamic simulations. In this paper, we systematically studied non-equilibrium ionization effects on clusters and the X-rays signatures using the same set of simulations we have developed \citep{WS09}. By using semi-analytic arguments together with numerical simulations, we have shown that the non-equilibrium ionization effect is nearly independent of cluster mass but depends strongly on redshift. In particular, non-equilibrium ionization effects are stronger for low-redshift clusters. Therefore, the brighter massive clusters at low-redshifts are good candidates for studying the non-equilibrium ionization effects. We systematically studied non-equilibrium ionization signatures in X-rays for a massive cluster with $M_{\rm sh} = 1.53 \times 10^{15}~M_{\odot}$. We first calculated the ionization fractions for 11 elements heavier than He following the electron temperature and density evolutions of each fluid element. We then calculated the X-ray emissivity of each fluid element and the resulting projected spectra for the cluster. Since the electron temperature profiles depend on electron heating efficiency $\beta$, we have considered three different possibilities which represent a very low heating efficiency ($\beta = 1/1800$), an intermediate heating efficiency ($\beta = 0.5$), and equipartition, $\beta = 1$. We also considered models which assume equilibrium ionization for comparison. At a radius (e.g., 2~Mpc) where the ionization timescale is long, the overall spectra for the NEI and CIE--Non-Eq models are very similar. This is because of the dominant free-free emission, and both models assume the same electron temperature. However, in the outer regions, e.g, at $r \sim 3.5$~Mpc which is between $R_{\rm vir}$ and $R_{\rm sh}$, the soft emission in the NEI model is dominated by line emission, where the CIE--Non-Eq spectrum is still dominated by the continuum free-free emission. By analyzing the surface brightness profiles, we found that soft emission (0.3--1.0~keV) for the NEI model can be enhanced by more than 20\% at around 3~Mpc, and up to nearly an order of magnitude near the shock radius compared to the CIE--Non-Eq model. The soft emission enhancement is mainly due to the line emission from under-ionized ions. The non-equilibrium ionization effects on the medium (1.0--2.0~keV) and hard (2.0--10.0~keV) band emissions are smaller. The overall X-ray band (0.3--10.0~keV) emission is dominated by the soft emission, and the total X-ray emission for the NEI model decrease much slower than that of the CIE--Non-Eq model. Thus, if cluster outer regions are in non-equilibrium ionization, the shock region will be much more luminous compared to the CIE--Non-Eq model. By inspecting a number of spectra, we found that the most prominent non-equilibrium ionization signature in line emission is the line ratio of the He-like \ion{O}{7} triplets and the H-like \ion{O}{8} doublets, $S($\ion{O}{8}$)/S($\ion{O}{7}$)$. The line ratios for the CIE models are higher than 10 for most regions between $r=$~1--4~Mpc, while the line ratios are smaller than 10 for the NEI models. The differences in the line ratios between the NEI and CIE models increase with radius, and the differences are more than an order of magnitude for radii beyond $\sim 2$~Mpc. These results are insensitive to the degree of non-equipartition or electron heating efficiency $\beta$. We suggest that the line ratios can be used to distinguish between the NEI and CIE models. The electron temperature profile can be determined from fits to the continuum spectra of the outer regions of clusters, allowing the CIE line ratios to be determined. Comparison to the observed ratios should show the effects of non-equilibrium ionization. Note that a line ratio of $S($\ion{O}{8}$) / S($\ion{O}{7}$) < 3$ in the outer region of a massive clusters is a clear signal of NEI. We have also studied the detectability of the \ion{O}{7} and \ion{O}{8} lines around cluster accretion shock regions with {\it IXO}, as well as the test for non-equilibrium ionization using the line ratio. For our optimum model, we found that with the XMS core array, an exposure time of 220~ksec is need to have a 3.0-$\sigma$ detection of the \ion{O}{7} lines and about 180~ksec is need to have $>30$ counts for a 3.2-$\sigma$ detection of the \ion{O}{8} lines. The uncertainties in NXB and GXB will not affect the results significantly. For the XMS full array while we assume the spectral resolution to be the same as the outer array throughout the detector, we found that the signal-to-noise ratios for our optimum model are higher for the same exposure time as the XMS core array. In particular, only about 130 (100)~ksec is needed to detect the \ion{O}{7} (\ion{O}{8}) line. The XMS full array is only slightly more subject to NXB and GXB uncertainties due to the poorer spectral resolution. To test the non-equilibrium ionization model without ambiguity requires measurements of both the electron temperature (or hardness ratio) and the line ratio so that the measured line ratio can be compared to the CIE line ratio inferred by the electron temperature. We have shown that this can be done within $\lesssim 2$~Mpc of a cluster by the {\it IXO} with sufficient confidence. Beyond $\sim 2.5$~Mpc where the surface brightness may be too low and measuring the electron temperature may be difficult, we have shown that if the line ratio is measured to be as low as $\sim 2$ at 3-$\sigma$ ($\Delta [S($\ion{O}{8}$)/S($\ion{O}{7}$)] \sim 2/3$ at 1-$\sigma$) in the outermost regions, this will rule out the CIE model at a 3-$\sigma$ level since the CIE line ratios are always higher than 4 for realistic cluster temperatures. A 3 or 4-$\sigma$ measurement of such a low line ratio is sufficient to provide a strong test of the non-equilibrium ionization. If the line ratio is measured to be even lower (e.g., $\lesssim 1$), only a 2 or 3-$\sigma$ will be sufficient to rule out the CIE model. On the other hand, if the line ratio is measured to be as high as 4, a 6-$\sigma$ measurement ($\Delta [S($\ion{O}{8}$)/S($\ion{O}{7}$)$] = 2/3 at 1-$\sigma$) will be necessary to rule out the NEI low line ratio of $\sim 2$ at a 3-$\sigma$ level. We found that an observation with about 130 (220)~ksec with the XMS full (core) array is enough to measure the line ratio at 2.3-$\sigma$. For a 3-$\sigma$ measurement of the line ratio, about 230 (380)~ksec will be needed for the XMS full (core) array, and this will provide a strong test for non-equilibrium ionization. In summary, detecting the \ion{O}{7} and \ion{O}{8} lines around the cluster accretion shock regions and testing non-equilibrium ionization in cluster outer regions with {\it IXO} are promising. It is expected that the \ion{O}{7} and \ion{O}{8} lines from WHIM will also be strong. Because of the high spectral resolution of the XMS, emissions from different redshifts should be easily separated. Only the emission from WHIM immediately surrounding the target cluster will be potentially confused with the emission from cluster outer regions. To observe the \ion{O}{7} and \ion{O}{8} lines and study NEI effects in cluster accretion shock regions, it will be best to avoid observing directions along the filaments where it is believed that denser preheated WHIMs and subclusters are preferentially accreted onto more massive and relaxed galaxy clusters. What do we learn about clusters from the ionization state of the outer gas? Since collisional ionization and recombination rates involve straightforward atomic physics, the processes are not in question and the rates are reasonably well-known. Unlike shock electron heating or rates for transport processes like thermal conduction, the basic physics is not uncertain and magnetic fields do not affect the results in a significant way. What we mainly learn about is the pre-shock physical state of materials which are being accreted by the cluster. If most of the WHIM is ionized beyond \ion{O}{7}, then the effects described in this paper will be greatly reduced. If most of the material currently being accreted by clusters comes in through filaments which have a higher ionization, then NEI effects will be diminished significantly. If most of the gas being added to clusters at present comes in through mergers with groups which deposit most of the gas in the inner regions of clusters, the gas will achieve CIE quickly. From the theoretical point of view, with the increasing number of observations of galaxy cluster outer regions ($\sim R_{200}$) and the potential to extend observations out to the shock radius with {\it IXO} in the future, it is necessary to perform more detailed simulations than ours. It is also interesting to extend our work to study the connections between the shocked ICM and the more diffuse WHIM surrounding clusters. Three-dimensional simulations will be essential to understand the effects of mergers or filament accretion on the degree of ionizations in different regions of clusters. This will allow us to characterize the variation of non-equilibrium signatures in the clusters; such calculations are essential to compare observational signatures with our understanding of the cluster physics near the accretion shocks. Cosmological simulations have been performed recently to study the NEI signatures \citep{YYS+03,CF06, YS06}. These studies have shown that both non-equipartition and non-equilibrium ionization effects are important in cluster outer regions; although they focus more on the lower density and lower temperature WHIM. High resolution simulations were also performed for studying NEI effects in clusters, but these are limited to binary mergers with idealized initial conditions and focus on the denser merger shocks \citep{AY08, AY10}. Re-simulating representative clusters and the surrounding WHIM from cosmological simulations with higher resolutions and including realistic physics (e.g., cooling, conduction, turbulent pressure, magnetic pressure, and relativistic support by cosmic rays) will be necessary to provide realistic model images and spectra. The different observational signatures and connections between the ICM and the more diffuse WHIM can also be addressed self-consistently by these simulations.

1012.0039

1012

1012.5336_arXiv.txt

{} {An analytical solution for the discrepancy between observed core-like profiles and predicted cusp profiles in dark matter halos is studied.} { We calculate the distribution function for Navarro-Frenk-White halos and extract energy from the distribution, taking into account the effects of baryonic physics processes.} {We show with a simple argument that we can reproduce the evolution of a cusp to a flat density profile by a decrease of the initial potential energy.} {}

There are many dwarf galaxies in the Universe (e.g.~\citet{mar97,saw10}), and the Local Group is an excellent place to study these systems and their properties \citep{mat98,pas10}. The nearly solid body rotation curve observed in dwarf galaxies indicates a central core in the dark matter distribution (e.g., \citet{bur95, deb02, gen05,deb05}). These results disagree with predictions from numerical simulations, which require density profiles with cusps (e.g. \citet{nav97,mor99,nav04}). \citet{nav96} were one of the first to point out the possibility that feedback mechanisms can turn the central dark-matter cusp into a cored one. Following this line of investigation, \citet{gneezao} analyzed the influence of winds. \citet{rea05} discussed how external impulsive mass-loss events can successfully act as a flattener for the central density cusps, and \citet{mas06} showed that random bulk motions of gas in small primordial galaxies could flatten the central dark matter cusp in $\sim 10^{8}$ years. Beyond that, several authors have suggested that the interstellar medium (ISM) of dwarf galaxies systems could be entirely removed as a result of star formation \citep{dek86,mor97,mur99,mac99,mor02,hen04, mor04}. In the present work, we construct a simple model under the hypothesis that astrophysical processes are able to remove baryonic gas from dwarf galaxies. We show for the first time a clear and analytical argument to solve the cusp problem. We argue that baryonic feedback could change the distribution function on the phase space of energy and angular momentum, flattening the density profile. This paper is organized as follows. In Sect. 2, we briefly discuss the removal of baryonic gas from dark matter halos as a consequence of different astrophysical processes. In Sect. 3, we review the formalism of distribution functions. In Sect. 4, we discuss the changes in density profile caused by gas removal and show our results. Finally we present our conclusions in Sect. 5.

The N-body simulations based on collisionless cold dark matter do not show the core-like behavior observed in local dwarf galaxies. Instead, they are better described by a steep power-law mass density distribution called cusp. Several attempts were made with numerical simulations in order to take into account the gas dynamics in the evolution of dark matter profiles. Considering that the total energy of the halos will change because of baryonic feedback, we showed that it is possible to model the evolution of an initial cusp-like profile into a core-like one. We used an analytical analysis based on the density profile DF. The advantage of this approach is that we obtain results similar to those extracted from numerical simulations, which take into account the gas dynamics \citep{pas10,gov09} in a very simple way. On top of that, considering the baryon loss as a mechanism of flattening the dark matter density profile of dwarf galaxies allows us to support the low amount of baryons observed in dwarf spheroidal galaxies \citep{gil2007}. However, as discussed in the introduction, there are two competitive processes occurring during the formation of galaxies, baryon accretion in the center of the halo will suffer adiabatic contraction and drag the dark matter profile into an even steeper one, and on the other hand, the baryonic characteristic of the gas causes it to radiate energy, increasing the entropy of the dark matter background and preventing the contraction. As our knowledge about these two processes is still incomplete, it is important to take into account these effects in future works.

1012.5336

1012

1012.2734_arXiv.txt

In general, for single field, the scale invariant spectrum of curvature perturbation can be given by either its constant mode or its increasing mode. We show that during slowly expanding or contracting, the spectrum of curvature perturbation given by its increasing mode can be scale invariant. The perturbation mode can be naturally extended out of horizon, and the amplitude of perturbation is consistent with the observations. We briefly discuss the implement of this scenario.

1012.2734

1012

1012.2672_arXiv.txt

We investigated the coronal activity of planet-hosting stars by means of statistical analysis for a complete sample of stars in the solar neighborhood. We find no observational evidence that Star-Planet Interactions are at work in this sample, at least not at the sensitivity levels of our observations. We additionally test the $\upsilon$~Andromedae system, an F8V star with a Hot Jupiter and two other known planets, for signatures of Star-Planet Interactions in the chromosphere, but only detect variability with the stellar rotation period.

Interactions between stars and close­-in planets can be expected from the analogy to binary stars. Binaries are often more active than single stars of the same spectral class \citep{AyresLinsky1980}, and X-­ray emission between the two components of a binary has been observed as well \citep{Siarkowski1996}. Thus, regarding stars with giant planets as binaries with an extremely small mass ratio, one expects to see enhanced activity levels of the host star from tidal or magnetic interaction with the planet \citep{CuntzSaar2000}, which should manifest themselves in activity proxies such as chromospheric Ca~II emission and coronal X-­ray emission. If Star­-Planet-­Interactions (SPI) are observed reliably, they can yield valuable information on the magnetic fields of exoplanets, the irradiation of exoplanetary atmospheres by the host star which in turn affects planetary evaporation \citep{Vidal-Madjar2003}, as well as orbital synchronization and planetary migration timescales.

Our investigations show that possible Star-Planet Interactions do not have a major influence on the average X-ray luminosity or $L_X/L_{bol}$ in nearby stars, at least not at the given sensitivity levels of our observations. Also our measurements of the chromospheric activity of the promising star-planet system $\upsilon$~And show no indications for SPI, but rather variability with the stellar rotation period. SPI seems to induce only small effects on the activity of the host stars; if observed over longer timescales and for more targets, however, they can provide insight into planetary and stellar magnetic fields.

1012.2672

1012

1012.1995_arXiv.txt

The eternal inflation scenario predicts that our observable universe resides inside a single bubble embedded in a vast inflating multiverse. We present the first observational tests of eternal inflation, performing a search for cosmological signatures of collisions with other bubble universes in cosmic microwave background data from the WMAP satellite. We conclude that the WMAP 7-year data do not warrant augmenting $\Lambda$CDM with bubble collisions, constraining the average number of detectable bubble collisions on the full sky $\nsavge < 1.6$ at $68 \%$ CL. Data from the {\em Planck} satellite can be used to more definitively test the bubble collision hypothesis.

1012.1995

1012

1012.3397_arXiv.txt

{} {The \object{Crab} nebula displayed a large $\gamma$-ray flare on September 18, 2010. To more closely understand the origin of this phenomenon, we analyze the INTEGRAL (20-500\,keV) and FERMI (0.1-300\,GeV) data collected almost simultaneously during the flare.} {We divide the available data into three different sets, corresponding to the pre-flare period, the flare, and the subsequent quiescence. For each period, we perform timing and spectral analyses to differentiate between the contributions of the pulsar and from the surrounding nebula to the $\gamma$-ray luminosity.} {No significant variations in the pulse profile and spectral characteristics are detected in the hard X-ray domain. In contrast, we identify three separate enhancements in the $\gamma$-ray flux lasting for about 12 hours and separated by an interval of about two days from each other. The spectral analysis shows that the flux enhancement, confined below $\sim$1\,GeV, can be modelled by a power-law with a high energy exponential cut-off, where either the cut-off energy or the model normalization increased by a factor of $\sim5$ relative to the pre-flare emission. We also confirm that the $\gamma$-ray flare is not pulsed. } {The timing and spectral analysis indicate that the $\gamma$-ray flare is due to synchrotron emission from a very compact Pevatron located in the region of interaction between the pulsar wind and the surrounding nebula. These are the highest electron energies ever measured in a cosmic accelerator. The spectral properties of the flare are interpreted in the framework of a relativistically moving emitter and/or a harder emitting electron population.}

With an integrated luminosity of about $5\times10^{38}\,\mathrm{erg\,s^{-1}}$ and a distance of $\sim$2\,kpc, the \object{Crab} supernova remnant is very bright from the radio domain to TeV energies \citep[see e.g.,][for a review]{2008ARA&A..46..127H}. It is powered by a pulsar spinning on its axis in about 33\,ms that injects energetic electrons into the surrounding nebula. Nearly all the nebular emission up to 0.4\,GeV is believed to be produced by synchrotron cooling of these electrons in an average magnetic field of $\sim300\,\mu$G. At higher energies, inverse Compton (IC) cooling dominates. The integrated high-energy flux of the nebula and the pulsar has been remarkably stable over the past few decades and the object is indeed used as a calibration source in several experiments \citep[but see ][ for the first report of a secular X-ray trend]{2010arXiv1010.2679W}. On 2010 September 22 the AGILE collaboration \citep{2010ATel.2855....1T} reported the first $\gamma$-ray flare from a source positionally consistent with the \object{Crab}, during which the flux above 100\,MeV was nearly double its normal value. The $\gamma-$ray flare from the direction of the \object{Crab} was confirmed by the Fermi collaboration \citep{2010ATel.2861....1B}. During a period partially covering the $\gamma$-ray flare, the \object{Crab} region was observed by the INTEGRAL satellite and the Swift/BAT telescope during its routine sky survey. No statistically significant increase in the \object{Crab} flux was observed in the IBIS/ISGRI light curves between 20 and 400 keV, as well as in the Swift/BAT $15-50\,$keV flux at a level of 5\% at the 1$\sigma$ confidence level \citep{2010ATel.2856....1F,2010ATel.2858....1M}. Swift/XRT observed the \object{Crab} region for 1\,ks on 2010 September 22 at 16:42 UT and did not reveal any significant variation in the source flux, spectrum, and pulse profile \citep{2010ATel.2866....1E}. The ultraviolet and soft X-ray images excluded the presence of any bright unknown field object that could have contributed to the $\gamma$-ray flux \citep{2010ATel.2868....1H}. The near-infrared \object{Crab} flux in the J and H bands was also constant \citep{2010ATel.2867....1K}. Dedicated pointing performed by RXTE did not show any significant change in the overall spectral properties in the 3-20\,keV band \citep{2010ATel.2872....1S}. A 5ks TOO Chandra observation was performed on 2010 September 28 to monitor the morphology of the inner nebula. There are no particular variations with respect to the over 35 previous observations, with the possible exception of an anomalous extension of a bright knot closer (down to 3'') to the pulsar, also noticed in one archival observation \citep{2010ATel.2882....1T}. A subsequent HST optical observation of the \object{Crab} confirmed an increase in the emission about 3 arcsec south-east of the pulsar with respect to archival observations \citep{2010ATel.2903....1C}. However, it remains unclear whether this feature is related to the $\gamma$-ray event. The \object{Crab} $\gamma$-ray flux returned to its usual level on 2010 September 23, less than a week after the onset of the flare \citep{2010ATel.2879....1H}.

To discriminate among the various models capable of reproducing the quasi-exponential turnover of the synchrotron emission of the nebula that peaks below the LAT energy window, we studied the Fermi data complemented by archival CGRO/COMPTEL data (0.75-30 MeV). A single power-law cannot reproduce the pre-flare data ($\chi^2 / d.o.f. \sim 44 / 15$). A power-law with a high energy exponential cutoff can instead reproduce the data ($\chi^2 / d.o.f. \sim 4 / 14$). To model the nebular synchrotron spectrum, we used the following function \begin{equation} \frac{\mathrm{d}N}{\mathrm{d}E} = N_0 \left( \frac{E}{1\,\mathrm{GeV}} \right)^{-\Gamma} \mathrm{exp}\left( - \frac{E}{E_{cutoff}} \right). \end{equation} The best-fit solution yields $\Gamma = -2.20 \pm 0.08$ and $N_0 = (4.3 \pm 1.9) \cdot 10^{-10}$ $\mathrm{ph~cm^{-2}~s^{-1}~MeV^{-1}}$. The difference between the quiescent and flaring spectra can be understood by considering two different extreme cases of either a constant power-law normalisation or a constant cutoff energy. In the former case, an increase in the energy cutoff of a factor of nearly 5 (from $77 \pm 15$ MeV to $367 \pm 45$ MeV) is needed (as illustrated in Fig.~\ref{fig:sed}). This increase is averaged over the whole flaring period, thus represents a lower limit, since in each single flare the cutoff energy might have been higher. In the latter case, the spectral variability can be explained by raising the continuum normalization by a factor of $\sim5$. \begin{figure}[t!] \resizebox{\hsize}{!}{\includegraphics[width=0.49\textwidth]{15980fg4.pdf}} \caption{\object{Crab} spectral energy distribution in the 100 MeV - 300 GeV energy range. The points with error bars are the Fermi detections before the flare (dark green), during the flare (red), and after the flare (light green). The black dashed line represents the contribution from the pulsar. The black dot-dot-dashed line represents the IC emission from the nebula. The blue and magenta dot-dashed and solid lines are the synchrotron nebula and the total emission before and during the flare, respectively. Arrows indicate the 95\% confidence flux limits.} \label{fig:sed} \end{figure} We note that the non-detection of any significant hard X-ray variability during the flare does not allow us to differentiate between the two possibilities as several electron populations are probably present in the nebula. As reported by \cite{1998nspt.conf..439A}, the COMPTEL data are characterized by a flattening of the spectrum that can be ascribed to the synchrotron emission of a separate electron population confined in compact regions such as wisps or knots. The luminosity of this component is less than 1\% of that of the whole nebula. The cutoff in the synchrotron spectrum occurs at a characteristic frequency $\nu_{\rm peak} \sim 4.2 \times 10^{6} \, B \, \gamma^{2}$. Provided that synchrotron radiation is the dominant mechanism through which particles channel their energy, the maximum electron Lorentz factor obtained by equating $t_{\rm sync}$ to $t_{\rm accel}$ is $\gamma_{\rm max} \propto (B \, \eta)^{-1/2}$, where $\eta \geq 1$ is the gyrofactor that characterizes the acceleration rate $\dot{\gamma}_{\rm accel} \equiv \gamma/t_{\rm accel}$ and $t_{\rm accel} = \eta \, E/q_{\rm e}\, B\,c$. This makes $\nu_{\rm peak}$ independent of $B$, leading to an electron synchrotron energy cutoff $\approx 160 \eta^{-1}$~MeV \citep[see e.g.][]{2000NewA....5..377A}. A higher value may imply that the conditions in the accelerator differ from those in the emission region, e.g. there is a lower magnetic field in the former, or that the synchrotron gamma-rays are produced in a relativistically moving region, which produces a shift in the energy cutoff to higher energies by the corresponding Doppler factor $\delta$. In this scenario, a value $\delta \sim 367/160 \approx 2.3$ would be required to explain the energy cutoff obtained during the flaring episode. On the other hand, magnetic fields at the level of between $\sim$~300~$\mu$G and $\sim$~2~mG are found in the synchrotron nebula and wisps, respectively \citep[see e.g.][]{2008ARA&A..46..127H}. Synchrotron radiation at $\sim$~1~GeV implies that $\gamma \sim 3-10 \times 10^{9}$ in the emitting regions. Taking $\delta \sim 2.3$, the comoving cooling timescale for those particles, taking an extreme value $B \sim$ 2 mG, is $\sim 0.3$~d. The corresponding observer timescale would then be similar to the decay time of the peaks present in the \textit{Fermi} lightcurve during the flaring period, $\lesssim 1$~d. In contrast, the flares could be related to an enhanced electron population, and the spectral variability could be obtained by raising the continuum normalization by a factor of $\sim 5$ or by adding a hard very high energy electron population (leading to a photon index $\Gamma<1$). In this case, a Doppler boosting may not be required, and the observed duration of the flares could correspond to the synchrotron timescale of PeV electrons embedded in magnetic fields $\lesssim 1\,\mathrm{mG}$. The duration of the three short flares limits the size of the emitting region(s) to $\lesssim 10^{15}$ cm. The peak luminosity of these flares is higher/brighter than $10^{35}$ erg/s, i.e. $\geqslant 0.5$\,\textperthousand \,\,of the \object{Crab} spin-down luminosity, assuming an isotropic distribution. The distance between the emitting region and the pulsar can thus be constrained to be $\leqslant 6\times 10^{16}$ cm, i.e. not larger than 15\% of the size of the bright synchrotron torus observed by Chandra and HST, and probably consistent with the half-width of this torus. The emitting region could therefore be linked to the interaction zone between the jet and the torus, which is found to have brightened in the HST image obtained on 2 October \citep{2010ATel.2903....1C}. The three flares separated by two days could possibly be related to various emitting knots in this region. Alternatively, gamma-rays could be produced within the jet itself. However, if the emitting region were moving at relativistic speeds, the emission would be radiated within an angle $\sim 1/\delta$. For reasonable values of the jet inclination angle with respect to the line of sight \citep[see e.g.][]{2004ApJ...601..479N}, this scenario would make the flares difficult to detect. To conclude, the flare relative short durations ($< 1$~day), their soft spectrum, and the analysis of the pulse profile in the 0.1-300\,GeV indicate that one or more compact portions ($\lesssim 10^{15}$~cm or $<$~0.1'') of the synchrotron nebula are responsible for the flares. In these region(s), freshly accelerated PeV electrons are rapidly cooling, causing the observed variability.

1012.3397

1012

1012.4781_arXiv.txt

As we approach solar convection simulations that seek to model the interaction of small-scale granulation and supergranulation and even larger scales of convection within the near-surface shear layer (NSSL), the treatment of the boundary conditions and minimization of sub-grid scale diffusive processes become increasingly crucial. We here assess changes in the dynamics and the energy flux balance of the flows established in rotating spherical shell segments that capture much of the NSSL with the Curved Spherical Segment (CSS) code using two different diffusion schemes. The CSS code is a new massively parallel modeling tool capable of simulating 3-D compressible MHD convection with a realistic solar stratification in rotating spherical shell segments.

} The solar differential rotation profile exhibits prominent radial shear layers near the top and bottom of the convection zone \citep{thomp03}. The near-surface shear layer (NSSL) occupies the upper 5\% of the Sun by radius, whereas the tachocline begins near the base of the convection zone. The dynamics of the NSSL are governed largely by vigorous granular-scale convection that is driven by radiative cooling and large thermodynamic gradients. The collective interaction of these granular-scale flows (average sizes of 1~Mm, lifetimes of 0.2~hr) is a major component in the formation of supergranular (15-35~Mm, 24~hr) and mesogranular (5-10~Mm, 5~hr) scales \citep{rast03,nord09}. Given this wide range of spatial and temporal scales, we currently cannot simultaneously model hundreds to thousands of supergranules, solar granulation, and deeper flows. However, we may still be able to characterize the influence of granulation and deep global flows on the NSSL by coupling CSS with global convection models from below \citep{miesch08} and with surface convection simulations from above \citep[e.g.][]{rempel09,nord09}. Such a coupling, whether it is statistical or direct, requires a careful treatment of diffusion and boundary conditions \citep{august10}. For instance, low diffusion is necessary to preserve the spatial structure and advective timescales of the downdrafts flowing into the CSS domain from the surface convection above. Given that our grid is five times coarser than that of typical surface convection simulations, this is no easy task. We have conducted two numerical simulations in a $20^\circ$ square patch centered on the equator that encompass most of the NSSL and some of the deep interior, rotating at the solar rate. These simulations explore the effects of turbulent-eddy and slope-limited diffusion schemes with an open lower radial boundary and closed upper boundary; as such they are identical except for the diffusion scheme. The governing equations and numerical approach used in solving them are briefly discussed in \S \ref{sect2}. The turbulent-eddy and slope-limited diffusion schemes are detailed in \S \ref{sect3}. The dynamics of the flows established in these simulations are examined in \S \ref{sect4}.

1012.4781

1012

1012.3732_arXiv.txt

Variability is a defining characteristic of young stellar systems, and optical variability has been heavily studied to select and characterize the photospheric properties of young stars. In recent years, multi-epoch observations sampling a wider range of wavelengths and time-scales have revealed a wealth of time-variable phenomena at work during the star formation process. This splinter session was convened to summarize recent progress in providing improved coverage and understanding of time-variable processes in young stars and circumstellar disks. We begin by summarizing results from several multi-epoch {\it Spitzer} campaigns, which have demonstrated that many young stellar objects evidence significant mid-IR variability. While some of these variations can be attributed to processes in the stellar photosphere, others appear to trace short time-scale changes in the circumstellar disk which can be successfully modeled with axisymmetric or non-axisymmetric structures. We also review recent studies probing variability at shorter wavelengths that provide evidence for high frequency pulsations associated with accretion outbursts, correlated optical/X-ray variability in Classical T Tauri stars, and magnetic reversals in young solar analogs.

The formation and early evolution of stars is inherently a time-domain process: the conversion of a dense molecular core into a zero-age main sequence star requires nearly every pertinent physical parameter (radius, density, temperature, v$_{rot}$, etc.) to change by multiple orders of magnitude over the relatively brief timescale of 10s of Myrs. These time-scales still dwarf that of a human lifetime\footnote{or, perhaps more relevantly, the timescale of a PhD thesis}, and one might {\it a priori} conclude that star formation is no more amenable to time-domain study than any other aspect of a stellar astrophysics. Photometric variability has nonetheless been recognized for decades as a common trait of many young stars \citep[e.g., ][]{Joy1945}. Historically, this variability has been best explored via optical photometry, most often sampling time-scales of days to (a few) years. These studies have provided detailed descriptions of the spot properties of optically revealed young stars \citep{Vrba1988}, a comprehensive inventory of stellar rotation in young clusters \citep[see first the review by][]{Herbst2007, Bouvier1995, Stassun1999, Rebull2004, Cieza2007, Irwin2008}, and a quantitative statistical portrait of accretion and extinction-induced variability \citep{Grankin2007}. These monitoring programs have also identified a number of rare, astrophysically valuable systems: pre-main sequence eclipsing binaries \citep[e.g., ][]{Cargile2008}, disk occulting systems \citep[ie, KH-15D:][]{Hamilton2001}, and stars undergoing massive accretion events \citep[i.e., FU Ori, EX Lup, or V1647-like variables:][]{Herbig1977, Herbig1989, Hartmann1996, Reipurth2004, Lorenzetti2007}. Significant advances in observational time domain astronomy are uncovering new phenomena and systems for study. This expansion is due to many factors: a steady improvement in the size and sensitivity of optical arrays, enhancing the coverage and cadence possible for observations of optically revealed clusters; even greater advances in the capabilities of near-infrared arrays; and access to precise multi-epoch mid-infrared (mid-IR) photometry and spectroscopy from the {\it Spitzer Space Telescope}. These new observational capabilities have allowed variability studies to cover a greater number of targets at higher cadences, and characterize the variability properties of more deeply embedded, and presumably less evolutionarily advanced, sources. The signature of these advances is evident in our increased sensitivity to variability at the lowest masses within nearby star-forming regions \citep{Cody2010} as well as the increasing frequency with which we identify formerly rare young variables, including new eclipsing binaries \citep{Hebb2010}, new disk occulting systems \citep[e.g., WL4:][]{Plavchan2008}, pulsating protostars \citep{Morales-Calderon2009} and large-amplitude outbursts \citep{Covey2010,Miller2010,Kospal2010,Garatti2010}. These observational advances have been accompanied by similar progress on the theoretical front, with increasingly detailed models of the physical processes underlying the observed variations \citep[e.g., accretion variations, disk processes, etc.;][]{Vorobyov2010, Zhu2010, Baraffe2010} This splinter session was convened to review recent progress in characterizing, analyzing, and understanding variability in the youngest stars. The session incorporated presentations covering physical processes that develop over a range of time-scales, with observational signatures spanning a wide range of wavelengths. Following these presentations, the audience participated in a broader discussion of these new results and highlighted areas of particular promise for future observational or theoretical work. We provide a brief summary of each presentation below, and conclude with a recap of the open questions highlighted in the audience discussion.

As this splinter summary demonstrates, recent years have seen considerable progress in our ability to observe, characterize, and model time-variable processes in the formation and early evolution of stars, circumstellar disks and planets. These advances include the ability to study variability across a wider range of wavelengths and time-scales than previously possible, and to construct more detailed and computationally intensive models of star and disk processes. Nonetheless, our current understanding of time-variable phenomena in early stellar evolution is significantly incomplete. In addition to the questions highlighted above, the audience identified several other outstanding challenges for studies of time-domain processes in the formation and early evolution of stars and planets. These include: \begin{itemize} \item{Improving the mechanisms and capabilities for conducting multi-site and multi-wavelength studies with long time baselines. } \item{Achieving statistically robust constraints on variability (e.g., FU Ori outburst rate/duty cycle), rather than qualitative descriptions.} \item{Identifying a critical mass of individual, rare systems (i.e, KH 15D, McNeil's Nebula, etc.) such that we can infer universal lessons from their specific properties.} \item{Developing better observational and theoretical diagnostics to locate the source regions and determine the causes of variability. Two examples are detecting kinematic signatures of the underlying gas motions in the line emission from the star and disk, and linking the photometric variability of the inner disk to changes in the scattered light from the spatially-resolved outer disk.} \item{Understanding the implications for planet formation of time-variable structure in protoplanetary disks.} \end{itemize} Meeting each of these challenges will require ingenuity, focused effort, and careful planning. The young stars presenting here today suggest that we will indeed be up to the task, but only time\footnote{(bad) Pun intended.} will tell.

1012.3732

1012

1012.0111_arXiv.txt

{It is currently accepted that intrinsically compact and bright radio sources characterized by a convex spectrum peaking at frequencies ranging from 100 MHz to a few GHz are young objects. Following the evolutionary models, these objects would evolve into the population of classical radio galaxies. However, the fraction of young radio sources in flux density-limited samples is much larger than what expected from the number counts of large radio sources. This may suggest that for some reason a significant fraction of young objects would never become large radio galaxies with sizes up to a few Mpc. The discovery of the young radio source PKS 1518+047 characterized by an uncommonly steep spectrum confirms that the radio emission may switch off shortly after its onset. Then the source spectrum steepens and evolves due to energy losses. If the interruption is not temporary, the fate of the fading sources is to disappear at frequencies lower than those explored by current radio telescopes. Fossils of past activities has been recently found at pc-scale distances from newly born radio sources, suggesting the presence of short-lived objects with an intermittent radio emission. } \FullConference{10th European VLBI Network Symposium and EVN Users Meeting: VLBI and the new generation of radio arrays\\ September 20-24, 2010\\ Manchester Uk} \begin{document}

Powerful (L$_{\rm 1.4 GHz}$ $>$ 10$^{25}$ W/Hz) and intrinsically compact ($<$ 1$^{\prime\prime}$) extragalactic radio sources represent a large fraction (15--30\%) of the radio sources selected in flux-limited catalogues. Their main characteristic is the steep synchrotron spectrum that turns over at frequencies between 100 MHz and a few GHz, and interpreted as due to synchrotron-self absorption \cite{mo08,snellen00}, although an additional contribution from free-free absorption (FFA) has been found in the most compact sources \cite{kameno00,marr01}. When observed with sub-arcsecond resolution these sources usually display a two-sided morphology with a weak core, jets and mini-lobes/hotspots, and for this reason they were termed compact symmetric objects (CSO) by \cite{wilkinson94}. Given their intrinsically compact size and their morphology resembling a scaled-down version of the classical powerful FRII \cite{fr} radio galaxies, CSOs have been interpreted as representing an early stage in the radio source evolution. Decisive support to this scenario came from the determination of both kinematic \cite{polatidis03} and radiative \cite{mm03} ages, resulting to be about 10$^{3}$--10$^{4}$ years, i.e. much smaller than the ages (10$^{7}$--10$^{8}$ years) estimated for classical radio galaxies with linear sizes up to a few Mpc \cite{lara00}. \\ In this context, it is possible to draw an evolutionary path in which CSOs are the precursors of extended radio galaxies \cite{phillips82}. Several evolutionary models \cite{fanti95,snellen00} have been developed aiming at describing how the physical properties, like luminosity and expansion velocity change as the radio source grows. However, many aspects, like the excess of young radio sources in flux-limited catalogues are not reproduced by the current models and additional explanations must be found.\\

1012.0111

1012

1012.5358_arXiv.txt

In this letter, we introduce a new method of image stacking to directly study the undetected but possible $\gamma$-ray point sources. Applying the method to the Australia Telescope 20 GHz Survey (AT20G) sources which have not been detected by Large Area Telescope (LAT) on {\it Fermi Gamma-ray Space Telescope} ({\it Fermi}) , we find that the sources contribute (10.5$\pm$1.1)\,\% and (4.3$\pm$0.9)\,\% of the extragalactic gamma-ray background (EGB) and have a very soft spectrum with the photon indexes of 3.09$\pm$0.23 and 2.61$\pm$0.26, in the 1--3 and 3--300\,GeV energy ranges. In the 0.1--1\,GeV range, they probably contribute more large faction to the EGB, but it is not quite sure. % It maybe not appropriate to assume that the undetected sources have the similar property to the detected sources.

The EGB was first detected by the SAS-2 mission \citep{fichtel75}, and its spectrum was measured with good accuracy by {\it Fermi} \citep[also called isotropic diffuse background,][]{abdo10a}. It has been found to be consistent with a featureless power law with a photon index of $\sim$2.4 in the 0.2--100\,GeV energy range and an integrated flux (E$\geq$100\,MeV) of 1.03$\times10^{-5}$\,ph cm$^{-2}$ s$^{-1}$ sr$^{-1}$. The origin of the EGB is one of the fundamental unsolved problems in astrophysics, and it has been a subject of study for a long time \citep[see][for a review]{k08}. The EGB could originate from either truly diffuse processes or from unresolved point sources. Truly diffuse emission can arise from numerous processes such as the annihilation of dark matter \citep{ahn07,c10,b10}, particle acceleration by intergalactic shocks produced during large scale structure formation \citep{gb03} etc. Blazars (including BL Lac objects, flat spectrum radio quasars, or unidentiffed flat spectrum radio sources) represent the most numerous population detected by the Energetic Gamma Ray Experiment Telescope (EGRET) on Compton Gamma Ray Observatory \citep{h99} and {\it Fermi} \citep{abdo10d}. Therefore, the blazars which have not been detected by the EGRET or LAT are the most likely candidates for the origin of the bulk of the EGB emission. Many authors have studied the luminosity function of blazars and showed that the contribution of blazars to the EGRET EGB could be in the range from 20\,\% to 100\,\% \citep{s96,n06,d07,c08,kn08,i09}. Nevertheless, starburst galaxy and non-blazar radio loud active galactic nuclei can also contribute a fairly large fraction of the EGB \citep{t07,bs09,b09}. Recently, \citet{abdo10b} built a source count distribution at GeV energy and yielded that point sources which had not been detected by the LAT can contribute 23\,\% of the EGB. At the fluxes currently reached by the LAT, they ruled out the hypothesis that point-like sources (i.e.\,blazars) produce a large fraction of the EGB. However, if the property of undetected sources is not similar to the detected sources, these conclusions maybe not correct. Therefore, we apply an image stacking method to directly study the undetected point sources. For a sample of possible $\gamma$-ray point sources which have not been detected by the {\it Fermi} due to their faint fluxes or soft spectra \citep{abdo10c}, we can stack a large number of them to improve the statistics \citep{ando10}. If their fluxes are not too faint, we can derive their mean flux and photon index by Maximum Likelihood (ML) method.

The stacked source is estimated to has a photon index of 2.81 and integrated flux of 1.07$\times10^{-7}$\,ph cm$^{-2}$ s$^{-1}$. % The TS is 129, corresponding to a significance of $\sim$11$\sigma$. The mean flux of these sources is 3.69$\times10^{-11}$\,ph cm$^{-2}$ s$^{-1}$, it is fainter than the faintest 1FGL source by a factor of 10. It can contribute 8.4\,\% of the EGB in the 1--300\,GeV energy range. We also apply our method to a subsample of flat spectrum radio sources (i.e. $\alpha_{\rm (5-20GHz)}<$0.5, with $F_{\nu}\propto\nu^{-\alpha}$, 1780 sources). Its photon index is 2.79, only slightly harder than the former. Its mean flux is 3.79$\times10^{-11}$\,ph cm$^{-2}$ s$^{-1}$, and the TS is 88. This subsample has not distinct characteristic from the other sources in $\gamma$-ray energy range. In order to test a more complicated spectral shape of the stacked source, we analyze the spectrum in the 1--10\,GeV energy range . We expect that the spectrum would be harder in this energy range. However, the estimated photon index is 3.01. Therefore, we analyze the spectrum in 2--10\,GeV, 3--10\,GeV energy range, respectively. The results are summarized in Table 1. We find that the spectrum is very soft in 1--3\,GeV energy range and becomes harder above 3\,GeV. It is indicated that two types of sources exist, in which one with softer and another with harder spectrum in GeV range. The former will dominate in lower energies, and latter in higher energies. Therefore, the spectrum of stacked source shows very soft in the 1--3\,GeV energy range. We will study this further if the optical properties of these objects can be obtained. Finally we obtain the properties of the spectrum as follows. In the 1--3\,GeV, the photon index is 3.12, the mean integrated flux is 3.89 $\times10^{-11}$\,ph cm$^{-2}$ s$^{-1}$, and the TS is 92. Obviously the flux is larger than that in the 1--300\,GeV energy range. It could be caused by that the spectrum in the 1--300\,GeV is not well fitted with a single power-law. In the 3--300\,GeV, the photon index is 2.66, the mean integrated flux is 3.72 $\times10^{-12}$\,ph cm$^{-2}$ s$^{-1}$, and the TS is 38. An decrease in $\ln{L}$ of 0.5 from its maximum value corresponds to the 68\% confidence (1 $\sigma$) region for each parameter \citep[see][]{m96}. We use this variance to estimate the error of each parameter. In three parameters ($\gamma_1$, $\gamma_2$ and $M$), we take two ones to be the values with maximal likelihood, and allow third one to change around its best value, we then test the deviation of $\ln{L}$ from its maximum value shown in figure 1. The 1 $\sigma$ errors of $\gamma_2$ are 0.25 and 0.22, the ones of $M$ are 270 and 53, in 1--3 and 3--300\,GeV respectively. The 1 $\sigma$ relative errors of fluxes are 10.5 \% and 17.3 \%, in 1--3 and 3--300\,GeV. In order to verify the effectiveness and accuracy of our method, we do the Monte Carlo (MC) simulations using the tool {\it gtobssim}. The simulating time is 26\,Ms, equaling to the time of real data we used. We simulate the Galactic and isotropic diffuse backgrounds using the models (e.~g. gll\_iem\_v02.fit, isotropic\_iem\_v02.txt) recommended by the LAT team, in which 3913 sources are generated each time, but only 2900 sources isotropically distribute on the sky with $|b|>15^\circ$. We complete one thousand MC simulations in 1--3 and 3--300\,GeV energy range using the obtained parameters. The diffuse source is simulated only once due to long run time, but its effect on the result is not remarkable because the source is random distribution and its photons are various. The distributions of the photon index, flux and TS for different energy ranges are shown in figure 2. They are compatible with Gaussian distributions. In the 1--3 and 3--300 \,GeV, the mean fluxes are 4.30 and 0.364 (in the unit of [$10^{-11}$\,ph cm$^{-2}$ s$^{-1}$]), their relative errors are 10.3\,\% and 20.7\,\%; the photon indexes are 3.15 and 2.71 with the errors of 0.23 and 0.26. The errors estimated here are similar to that found in the fourth paragraph. Comparing the input parameters, we find that the systematic errors occur, especially for the flux in the 1--3\,GeV energy range. They could be caused by that diffuse background source can not be described by a single power-law spectrum. Because the MC method can obtain the systematic errors, we use this method to correct our results as follows: in the 1--3 and 3--300\,GeV, the fluxes are 3.48$\pm$0.36 and 0.380$\pm$0.080 (in unit of [$10^{-11}$\,ph cm$^{-2}$ s$^{-1}$]), and the photon indexes are 3.09$\pm$0.23 and 2.61$\pm$ 0.26 respectively, while the contribution to the EGB is (10.5$\pm$1.1)\,\% and (4.3$\pm$0.9)\,\%, which are much smaller than the result (17\,\%) of \citet{G10b}. If the soft spectrum in 1--3\,GeV is caused by the spectral broken of some sources, the photon index would not be extrapolated to lower energy range. However, as long as the spectrum of stacked source is not harder than the EGB, the contribution to the EGB in 0.1--1\,GeV will be larger than that in 1--3\,GeV. Our result is compatible with the result (23\,\%) of \citet{abdo10b} because other point sources could contribute to the EGB. In this letter, we introduce a new method of images stacking to directly study the contribution of undetected point sources to the EGB. Our method is more direct than the methods used by many authors. Those methods involve the $\gamma$-ray luminosity of undetected sources which is estimated through the properties of a few detected sources. They include many uncertainties and lead the result to be questionable validity. Applying our method, we find that the undetected sources in AT20G can contribute (10.5$\pm$1.1)\,\% and (4.3$\pm$0.9)\,\% to the EGB in the 1--3 and 3--300\,GeV energy range respectively. Their $\gamma$-ray spectrum is very soft, implying that the emissive property is different for undetected and detected sources. Applying our method to estimate the contribution of all point sources to the EGB, we need a complete sample of possible $\gamma$-ray point sources which is not easily constructed. In this letter, we only estimate the contribution of AG20G to the EGB. We will study more samples of possible $\gamma$-ray point sources in the future.

1012.5358

1012

1012.5077.txt

We use the Mars Regional Atmospheric Modeling System (MRAMS) to simulate lake storms on Mars, finding that intense localized precipitation will occur for lake size $\geq$10$^3$ km$^2$. Mars has a low-density atmosphere, so deep convection can be triggered by small amounts of latent heat release. In our reference simulation, the buoyant plume lifts vapor above condensation level, forming a 20km-high optically-thick cloud. Ice grains grow to 200 $\mu$m radius and fall near (or in) the lake at mean rates up to 1.5 mm/hr water equivalent (maximum rates up to 6 mm/hr water equivalent). Because atmospheric temperatures outside the surface layer are always well below 273K, supersaturation and condensation begin at low altitudes above lakes on Mars. In contrast to Earth lake-effect storms, lake storms on Mars involve continuous precipitation, and their vertical velocities and plume heights exceed those of tropical thunderstorms on Earth. For lake sizes 10$^{2.5}$ - 10$^{3.5}$ km, plume vertical velocity scales linearly with lake area. Convection does not reach above the planetary boundary layer for lakes $\ll$10$^3$ km$^2$ or for atmospheric pressure $>$ O(10$^2$) mbar. Instead, vapor is advected downwind with little cloud formation. Precipitation occurs as snow, and the daytime radiative forcing at the land surface due to plume vapor and storm clouds is too small to melt snow directly ($<$ +10 W/m$^2$). However, if orbital conditions are favorable, then the snow may be seasonally unstable to melting and produce runoff to form channels. We calculate the probability of melting by running thermal models over all possible orbital conditions and weighting their outcomes by probabilities given by long-term integrations of the chaotic diffusion of solar system orbital elements \citep{las04}. With this approach, we determine that for an equatorial vapor source, sunlight 15\% fainter than at present, and snowpack with albedo 0.28 (0.35), melting may occur with 4\% (0.1\%) probability. This rises to 56\% (12\%) if the ancient greenhouse effect was modestly (6K) greater than today. %Although $>$ 50\% of vapor released lands near the vapor source as snow, a significant amount escapes the boundaries of our simulation and so is available to drive transient global climate change.

Evidence for runoff on Mars shows it to be patchy in both space and time \citep{kra08,wil07,wei08,fas08b,hyn10,car00}, so perhaps past precipitation was also patchy. Because patchy surface vapor sources cannot persist in equilibrium with a dry atmosphere \citep{ric08,sot08}, vapor would have to be supplied from an environment not in equilibrium with surface conditions. Such environments can be transient, such as an impact lake, or long-lived, such as the base of a wet-based ice-sheet. They can be high-temperature, such as fumaroles (or a lava flow advancing over snowpack), or involve only moderate temperatures, such as groundwater discharge. Here we use a mesoscale model to explore the atmospheric response to one example of a non-equilibrium vapor source: an ephemeral lake on a cold desert Mars. We track the fate of vapor supplied by the lake from release, through cloud formation, to precipitation, and consider whether the resulting snow will melt and provide runoff to form channels. Lake size, solar luminosity, lake geometry, and atmospheric pressure all affect the results. Only idealized results are presented: a companion paper (\citet{kit102}; henceforth Paper 2) uses the same model for a case study of the Juventae plateau inverted channel networks \citep{wei08}. Low volumetric heat capacity makes the Mars atmosphere's response to lake vapor release similar to tropical moist convection on Earth, so we borrow ideas from tropical meteorology to understand our results (e.g. \citet{ema94}). Figure \ref{EQUIVTEMP} shows the low-pressure lake effect: condensation of a small amount of vapor in a thin atmosphere can produce strong convection, which in Earth's thick atmosphere would require condensation of a large amount of vapor and correspondingly high water surface temperatures. Localized precipitation on a cold desert planet is normally transient precipitation. A warm, wet patch connected to the global atmosphere will lose water to cold traps elsewhere on the planet \citep{ric08}; the water table will withdraw to the subsurface because of evaporative losses \citep{sot08}. In the absence of an external heat source, evaporative and radiative cooling will cause any lake to quickly freeze \citep{lor05,con10}. (These arguments do not apply to springs, nor proglacial discharge of subglacial meltwater. In these cases, under cold conditions, any given parcel of water will freeze over, but a sustained vapor source can nevertheless exist at the discharge site.) As the ice thickens, ice surface temperature will fall and the lower saturation vapor pressure will cause the rate of vapor release to greatly decrease. For realistic external heat sources, the lake lifetime is still short. For example, consider an impact-generated lake near the freezing point overlying shocked basalt that is initially at 1000 $^\circ$C. The lake is assumed to be well-mixed by waves driven by lake-effect and impact-thermal storms, and convection driven by bottom heating. Icing-over is inevitable when the heat flow from the interior of the lake toward the surface is less than the evaporative and radiative losses at the surface. The maximum time before icing over, $t$, is therefore \begin{equation} t \approx \frac{D (T_{b} - T) c_{b} \rho_{b} }{ (E L_{vap} + \sigma T^4)} \end{equation} %t = 100 * 800* 800 * 3000 / (1200 + 300) where $D$ is the depth of pervasive fracturing within the rock ejecta, $T_b$ = 1273K is the initial temperature of the basalt, $T$ = 278K is lake surface temperature, $c_{b}$ = 840 J/kg/K the specific heat capacity of the basalt, $\rho_{b}$ = 2000 kg/m$^3$ the density of the fractured basalt, $E$ = 2 kg/m$^2$/hr the evaporation rate, $L_{vap}$ = 2.5 x 10$^5$ J/kg the latent heat of vaporisation, and $\sigma$ = 5.67 x 10$^{-8}$ W/m$^2$/K$^4$ is the Stefan-Boltzmann constant. For $D$ = 100m we obtain $\sim$4 Earth years: a geological instant. The true timescale will be less. For example, if fractures within the ejecta anneal, the relevant timescale is conductive cooling of a half-space (the ejecta layer) by an isothermal boundary condition (the well-mixed lake) until the heat flow into the bottom of the lake is less than heat loss at the top of the lake \citep{tur02}: \begin{equation} t \approx \left( \frac{k (T_{b} - T) }{ E L_{vap} + \sigma T^4 } \right)^2 \frac{1}{ (\pi \kappa)} %t \approx \left( \frac{(E L_{vap} + \sigma T^4) \sqrt{\kappa \pi}}{2 k (T_{b} - T)} \right)^2 \end{equation} In this case, the conductive heat flow can only balance evaporative plus radiative losses for $\sim$3 days for thermal diffusivity $\kappa$ = 10$^{-6}$ m$^2$ s$^{-1}$: after this, an ice cover must form. Therefore, we are interested in spatially restricted (10$^0$ - 10$^3$ km) water sources which cease to emit vapor in timescales $<$ 1 year. This is the domain of mesoscale modelling. We use the Mars Regional Atmospheric Modeling System (MRAMS), also used for entry, descent and landing simulations for the Mars Exploration Rovers, Mars Phoenix, and Mars Science Laboratory (Appendix A; \citet{raf01,mic08}). MRAMS explicitly resolves the size spectrum of dust and water ice aerosol for both cloud microphysics and radiative transfer, so it is well-suited for our cloud-forming numerical experiments \citep{mic08}.

We conclude from this study that:- (1) A low-pressure lake effect creates deep convection, rapid updrafts, and intense precipitation above cold liquid water surfaces on Mars. On Earth, because of much higher atmospheric pressure, this requires tropical temperatures. (2) We use MRAMS to simulate lake storms on Mars. The modeled storms have updraft velocities and plume heights which exceed the intensity of the strongest recorded thunderstorms on Earth. (3) The fraction of vapor that is trapped near the vapor source as localized precipitation increases with lake size. (4) The localized greenhouse effect of the released water vapor is too weak to cause melting of the snow. (5) Melting of equatorial, rapidly-emplaced localized snow with the albedo of dust is possible for a subset of orbital conditions, even with the present-day weak greenhouse effect. (6) Assuming that transient lakes on Mars are uncorrelated with orbital forcing, melting of rapidly-emplaced localized precipitation is more likely than melting of precipitation that has been emplaced in equilibrium with orbital conditions. (7) Taken together, localized storms, rapidly-emplaced localized precipitation, and favorable orbital conditions provide an alternative working hypothesis for (at least part of) the erosion and channel formation observed on pre-modern Mars. %%% End of body of article: %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %% Optional Appendix goes here % %%%%%%%%%%%%%%%%% % Geophysical Research Letters only allows an appendix without a letter. %% You can get this result with %

1012.5077

1012

1012.2052_arXiv.txt

We report the discovery of the first \ion{He}{1}*$\lambda 10830$ broad absorption line quasar FBQS~J1151$+$3822. Using new infrared and optical spectra, as well as the SDSS spectrum, we extracted the apparent optical depth profiles as a function of velocity of the 3889\AA\/ and 10830\AA\/ \ion{He}{1}* absorption lines. Since these lines have the same lower levels, inhomogeneous absorption models could be used to extract the average true \ion{He}{1}* column density; the log of that number was 14.9. The total hydrogen column density was obtained using {\it Cloudy} models. A range of ionization parameters and densities were allowed, with the lower limit on the ionization parameter of $\log U=-1.4$ determined by the requirement that there be sufficient \ion{He}{1}*, and the upper limit on the density of $\log n = 8$ determined by the lack of Balmer absorption. Simulated UV spectra showed that the ionization parameter could be further constrained in principle using a combination of low and high ionization lines (such as \ion{Mg}{2} and \ion{P}{5}), but the only density-sensitive line predicted to be observable and not significantly blended was \ion{C}{3}$\lambda 1176$. We estimated the outflow rate and kinetic energy, finding them to be consistent but on the high side compared with analysis of other objects. Assuming that radiative line driving is the responsible acceleration mechanism, a force multipler model was constructed. A dynamical argument using the model results strongly constrained the density to be $\log n \geq \sim 7$. Consequently, the log hydrogen column density is constrained to be between 21.7 and 22.9, the mass outflow rate to be between 11 and 56 solar masses per year, the ratio of the mass outflow rate to the accretion rate to be between 1.2 and 5.8, and the kinetic energy to be between 1 and $5\times 10^{44}\rm \, erg\, s^{-1}$. We discuss the advantages of using \ion{He}{1}* to detect high-column-density BALQSOs and and measure their properties. We find that the large $\lambda f_{ik}$ ratio of 23.3 between the 10830\AA\/ and 3889\AA\/ components makes \ion{He}{1}* analysis sensitive to a large range of high column densities. We discuss the prospects for finding other \ion{He}{1}*$\lambda 10830$ BALQSOs and examine the advantages of studying the properties of a sample identified using \ion{He}{1}*.

Active galactic nuclei (AGN) are powered by mass accretion onto a supermassive black hole. But while most of the gas is accreted by the black hole, some fraction of the gas is blown out of the central engine in powerful winds. These outflows are seen as the blue-shifted absorption lines primarily in the rest-frame UV spectra for about $\sim 50$\,\%\ of AGN as narrow absorption lines \citep{crenshaw03} and for about $\sim 10$ to 40\,\%\ as broad blue-shifted absorption troughs \citep[e.g.,][]{weymann91, gibson09, dai08}. Outflows are an essential part of the AGN phenomenon because they can carry away angular momentum and thus facilitate accretion through a disk. Winds are important probes of the chemical abundances in AGN, which appear to be elevated \citep{hf99}. They can distribute chemically-enriched gas through the intergalactic medium \citep{cavaliere02}. They may carry kinetic energy to the host galaxy, influencing its evolution, and contributing to the coevolution of black holes and galaxies \citep[e.g.,][]{so04}. Physical properties, including the mass and origin of AGN outflows, remain largely mysterious. It is particularly important to understand the acceleration mechanism. Resonance-line absorption (i.e., continuum photons are absorbed as permitted transitions in ions) is a promising mechanism, and there is compelling evidence that it is present in a few objects \citep[e.g.][]{arav95,north06}. Other mechanisms include hydromagnetic acceleration \citep{everett05} and acceleration due to dust \citep{sn95}. As discussed by \citet{ge07}, different acceleration mechanisms may dominate in different parts of the outflow. These models could in principle be distinguished by measuring the properties of the outflow. So, for example, while resonance-line absorption is a compelling mechanism (since we see the troughs created by absorption), there is evidence in some cases that the mass-outflow rates may be too high for it to be feasible (e.g., Hamann et al.\ 1997; Leighly et al.\ 2009). Thus, kinetic energy is a sensible discriminant for these models; e.g., we ask, is there sufficient energy available to accelerate wind of a particular column density to a particular velocity? Determining the kinetic energy deposited into the wind requires measurement of the velocity, which can be obtained directly from the absorption line profile, assuming that the flow is radial, and the column density, which is much more difficult to measure. The UV lines have apparent optical depths generally less than one, so it appears that column densities could be determined by simply integrating over the apparent optical depth profile. However, it is now known that the lines are generally saturated, although not black, implying that the absorbing material only partially covers the source \citep[e.g.,][]{sabra05}. The column density and covering fraction can be solved for in special cases when two or more lines from the same lower level can be measured \citep[e.g.,][]{hamann97, arav05}. For example, atomic physics requires that the 1548\,\AA\ component of the \ion{C}{4} doublet have twice the optical depth of the 1551\,\AA\ component. Partial covering gives an apparent optical depth ratio of less than two. The measurements of the two lines can be used to solve for the two unknowns, the optical depth and covering fraction. This method, while powerful, has limitations. Blending can be a problem when the lines are broad. For example, the \ion{C}{4} doublet at 1548 and 1551\AA\/ is a suitable pair of lines for this method, since these two transitions have the same lower level. However, their separation is only 2.6\AA\/, corresponding to only $500\rm \, km\, s^{-1}$. While these lines may be resolved in narrow-line objects, they will be profoundly blended in objects with broad lines. In addition, this method fails if the lines are saturated. Lines from permitted transitions in ions from relatively abundant elements saturate easily at relatively low column densities. Examples of such lines are those that are most easily recognized in BALQSOs, such as \ion{C}{4}. \ion{P}{5}$\lambda\lambda 1118, 1128$ has been recognized as a valuable probe of higher column densities \citep{hamann98} because of its low elemental abundance, only $9.3 \times 10^{-4}$ that of carbon \citep{grevesse07}. So, from the observation of \ion{P}{5} in a quasar, we can generally infer that \ion{C}{4} and similar lines are saturated, even though they may not be black. However, \ion{P}{5} has the problem in that it falls in the far UV part of the spectrum. It is therefore, in practice, accessible only by space-based observatories, except in BALQSOs with redshifts greater than $\sim 2.5$. In high-redshift objects, however, Ly$\alpha$ forest lines ablate the spectrum in the far UV, and \ion{P}{5} may be difficult to measure. In this paper, we report the discovery of the first \ion{He}{1}*$\lambda 10830$ broad absorption line quasar FBQS~J1151$+$3822. Absorption has been seen in the \ion{He}{1}*$\lambda 3889$ transition (a transition to 3p from the metastable 2s level in the \ion{He}{1} triplet) in several quasars and Seyfert galaxies \citep[e.g., Mrk~231;][]{boksenberg77}, but it has never before been reported in the 10830\AA\/ transition (a transition to 2p from the metastable state). We also discuss the use of \ion{He}{1}* absorption lines at 3889 and 10830\AA\/ for detecting low-redshift, high-column-density BALQSOs. \ion{He}{1}* offers a number of advantages. The metastable state is populated by recombination from He$^+$, so it is a high-ionization line, produced in the same gas as the usual resonance lines such as \ion{C}{4}$\lambda 1549$. The metastable state has a low abundance, comparable to that of $P^{+4}$. The two \ion{He}{1}* transitions are widely separated and located in a region of the spectrum where there are few other absorption lines, so blending is not a problem. They are located in the optical and infrared, so observation is possible from the ground. Observation is then limited to lower-redshift quasars, but turns out to be a valuable property because there are now only a handful of low-redshift BALQSOs known, many discovered serendipitously in space-UV spectra taken for other reasons. Finally, these two lines have a high ratio of $\lambda f_{ik}=23.3$, compared with $\sim 2$ for the usual resonance lines. This means that the two lines together are useful over a larger range of high column densities. Although the \ion{He}{1}*$\lambda 10830$ line has so far not been used extensively in the study of quasars, its value has been recognized for the study of other objects, including winds in young stars \citep[e.g.,][]{edwards03}. How metastable is the \ion{He}{1}* 2s triplet state? The calculation of the rate of decay for this and for \ion{H}{1}* 2s has a fascinating history. Electric dipole transitions are strictly forbidden. Electric quadrupole and magnetic dipole transitions are forbidden. However, decay via a two-photon process is possible \citep{bs57}. The rate for the two-photon process was computed for \ion{H}{1} and discussed for \ion{He}{1} by \citet{bt40}. They found that while the lifetime of the singlet state of \ion{He}{1} should be similar to that of hydrogen, at $0.11$--$0.15$ seconds, the lifetime of the triplet state should be about $10^{-6}$ times smaller, or $\sim 1.5$ days. The lifetime was then calculated by \citet{m57} to be about 0.5 days. But \citet{dd68} showed that the analysis performed by \citet{bt40} was not applicable to the triplet state, because while the singlet state decay has net $\Delta m=0$, the triplet state decay transports one quantum of angular momentum. In \citet{dvd69} they show that the lifetime of the metastable state for the two-photon decay should be 7.9 years. But then features from transitions from the triplet state to the ground from helium-like ions were observed in the laboratory and the sun, implying that there must be a direct transition from the metastable state to the ground \citep{gj69}. The key lay apparently in the relativistic corrections to the magnetic dipole transition. \citet{griem69} derived rates commensurate with the observation using an approximate Dirac theory. Finally, \citet{drake71} calculated the rate using higher order terms yielding the current best estimate of then lifetime of 2.2 hours. This paper is organized as follows. In \S 2, we describe the infrared and optical observations of FBQS~J1151$+$3822, the extraction of the apparent optical depth profiles, and application of partial covering models to extract the true column density and covering fraction. In \S 3, we use {\it Cloudy} modeling to extract column density information from the results, and to compute synthetic UV spectra. \S 4 provides a discussion of the kinetic luminosity, the viability of acceleration mechanisms, the column densities and covering fractions that can be reliably measured using the \ion{He}{1}* lines, and the prospects for using \ion{He}{1}* to find low-redshift quasars. A brief appendix explores the dependence of our total hydrogen column density estimates on the spectral energy distribution, the metallicity, and turbulence. We use cosmological parameters $\Omega_\lambda=0.73$, $\Omega_{M}=0.27$,$H_0=71\rm\, km\,s^{-1} Mpc^{-1}$ unless otherwise specified.

The principal results and findings of this paper are as follows. \begin{itemize} \item We report the first short wavelength (0.8--2.4$\mu$) infrared spectroscopic observation of FBQS~J1151$+$3822 using SpeX on the IRTF, as well as a new optical spectroscopic observation of FBQS~J1151$+$3822 using CCDS on the 2.4 meter Hiltner telescope at MDM observatory. In addition, we analyzed the SDSS spetrum. We discovered broad \ion{He}{1}* absorption lines, both in the 3889\AA\/ transition and, for the first time in a quasar, in the 10830\AA\/ transition. The terminal velocity $v_{max}$ is $11,000\,\rm km\, s^{-1}$, indicating that FBQS~J1151$+$3822 is a broad absorption line quasar. \item We extracted the apparent optical depths for the absorption lines. The procedure was straightforward for the 10830\AA\/ line because the IR continuum is fairly simple. But for the 3889\AA\/ line, this procedure was complicated by strong \ion{Fe}{2} emission, so a template-fitting method had to be used. An upper limit on Balmer absorption was extracted from the SDSS spectrum. \item The \ion{He}{1}* column densities were first estimated by integrating directly over the optical depth profiles. The estimate based on the 3889\AA\/ line ($\log N_{HeI*} \sim 14.85$) was substantially higher than the one based on the 10830\AA\/ line ($\sim 14.3$). This result implies that the absorber does not fully cover the source. Since the two transitions have the same lower (metastable) level, we can solve two-parameter inhomogeneous covering models as a function of velocity. We use a pure partial covering model and a power law model. Both models give the same average log \ion{He}{1}* column of 14.9. The hydrogen $n=2$ limit depends on the details of the population of the $2p$ and the $2s$ states; assuming that the $2s$ metastable state was solely populated yielded an upper limit of log hydrogen $n=2$ column of 13.5. \item {\it Cloudy} modeling was used to constrain the range of parameter space permitted by the measurements. The ionization parameter $\log U$ must be greater than or equal to $-1.4$ in order to produce sufficient \ion{He}{1}*. A log density less than $\sim 8$ was required, or too much Balmer absorption was predicted. The log hydrogen column density was constrained to be greater or equal to 21.6. Comparison with other objects showed that FBQS~J1151$+$3822 is a relatively high-ionization, high-column-density BALQSO. \item The {\it Cloudy} simulations were used to see if UV spectra would be able to constrain parameters further. It was found that ionization parameter would be reasonably easily constrained through observation of low ionization lines such as \ion{Mg}{2} and \ion{N}{3}, which decrease in optical depth as the ionization parameter increases, and observation of the high ionization line \ion{P}{5}, which increases in optical depth as the ionization parameter increases. Several density diagnostic lines were examined; all were predicted to be too blended to be useful except for the high-velocity side of \ion{C}{3}*$\lambda 1176$. \item The broad band photometry was examined. The observed decrease in the UV was too steep to be solely caused by reddening. Attenuation is expected also because of the absorption lines predicted in the UV. However, a combination of correction for modest reddening ($E(B-V)=0.1$ for an SMC reddening curve) and attenuation due to predicted high-ionization absorption lines left the observed UV still a factor of $\sim 2.5$ too low. \item The mass outflow rate, kinetic luminosity and ratio of kinetic to bolometric luminosity were estimated. The range was large over allowed parameter space but generally the minimum values were higher than those estimated in other objects. \item Acceleration mechanisms were examined. A force-multiplier analysis was done using the {\it Cloudy} results. It was found for this luminous object, radiating at about 1/2 the Eddington luminosity, the terminal velocity $v_{max}$ could be easily achieved. Furthermore, a comparison of the size scale of the outflow with the radius of the outflow constrained the log of the density to be greater than $\sim 7$. This strongly constrained parameter space, yielding an acceptable range log hydrogen column densities of 21.7--22.9, radii between 2--12 parsecs, mass outflow rates between 11 and 54 solar masses per year, ratio of outflow to inflow rates between 1.2 and 5.8, kinetic luminosity between 1 and $5\times 10^{44}\rm \, erg\, s^{-1}$, and kinetic luminosity between 0.2 and 0.9\% of the bolometric luminosity. Compared with other objects, FBQS~J1151$+$3822 has a relatively powerful outflow originating relatively close to the central engine. \item We examine the potential for using the \ion{He}{1}* absorption lines to detect and measure the properties of high-column-density BALQSOs. Objects with high column densities may require the highest kinetic energies and so may be important for testing models. We demonstrate that many prominent BAL lines become saturated at high column densities. We show that \ion{He}{1}*$\lambda 10830$ compares favorably with \ion{P}{5}, a widely used probe of high column densities, while \ion{He}{1}*$\lambda 3889$, because of its low oscillator strength, is even more sensitive. We use simulated spectra to see how well we can reproduce input covering fractions and column densities using a partial covering model. We find that the high $\lambda f_{ik}$ ratio of 23.3 makes the \ion{He}{1}* lines more sensitive to a wider range of column density than lines with a $\lambda f_{ik}$ ratio equal to 2 (like \ion{P}{5}). \item We briefly discussed the prospects of finding other \ion{He}{1}*$\lambda 10830$ BALQSOs. We have in fact already collected data on six additional objects; furthermore, we have discovered that several well-known, bright low-redshift BALQSOs have no \ion{He}{1}*$\lambda 10830$ absorption, a fact that will place upper limits on the column densities in those objects. We discussed the utility of observing a sample of low-redshift BALQSOs identified via \ion{He}{1}* absorption. Such objects may be valuable for understanding the relationship between luminosity and terminal velocity $v_{max}$ by allowing definition of a lower luminosity sample in a region currently sparsely populated. \end{itemize}

1012.2052

1012

1012.4168_arXiv.txt

We present spectroscopic follow-up observations of 70 $\mu$m selected galaxies from the SWIRE XMMLSS and Lockman Hole fields. We have measured spectroscopic redshifts for 293 new sources down to a 70 $\mu$m flux limit of 9mJy and $r<$ 22 mag. The redshift distribution peaks at $z\sim$ 0.3 and has a high redshift tail out to $z=3.5$. We perform emission line diagnostics for 91 sources where [OIII], H$\beta$, [NII], H$\alpha$ and [SII] emission lines are available to determine their power source. We find in our sample 13 QSOs, 1 Seyfert II galaxy, 33 star forming galaxies, 30 composite galaxies, 5 LINERs and 21 ambiguous galaxies. We fit single temperature dust spectral energy distributions (SEDs) to 81 70 $\mu$m sources with 160 $\mu$m photometry to estimate dust temperatures and masses. Assuming the dust emissivity factor ($\beta$) as 1.5, we determine dust temperatures in the range $\sim$ 20-60K and dust masses with a range of 10$^6$-10$^9$ M$_\odot$. Plotting these objects in the luminosity-temperature diagram suggests that these objects have lower dust temperatures than local IR luminous galaxies. The \textit{Herschel} Space Observatory will be crucial in understanding the nature of these sources and to accurately determining the shape of the Rayleigh-Jeans tail of the dust SED. We then model SEDs from optical to far-IR for each source using a set of galaxy and quasar templates in the optical and near-IR (NIR) and with a set of dust emission templates (cirrus, M82 starburst, Arp 220 starburst and AGN dust torus) in the mid-IR (MIR) to far-IR (FIR). The number of objects fit with each dust template are: 57 Arp 220, 127 M82, 9 cirrus, 1 AGN dust torus, 70 M82 and cirrus, 26 M82 and AGN dust torus and 3 Arp 220 and AGN dust torus. We determine the total IR luminosity (L$_\mathrm{IR}$) in range 10$^8$-10$^{15}$ L$_\odot$ by integrating the SED models from 8 to 1000 $\mu$m.

The discovery of the Cosmic Infrared Background (CIRB; \citealt{puget96, fixsen98} has shown that at least half of the energy generated by star-formation and Active Galactic Nuclei (AGN) in the Universe has been absorbed by dust and re-radiated in the IR \citep{gispert00,hauser01}. A significant fraction of the background has been resolved into a population of IR sources, luminous (LIRG: L$_\mathrm{IR}$ = 10$^{11}$-10$^{12}$ L$_\odot$), ultraluminous (ULIRG: L$_\mathrm{IR}$ = 10$^{12}$-10$^{13}$ L$_\odot$) and hyperluminous (HLIRG: L$_\mathrm{IR} >$ 10$^{13}$ L$_\odot$) infrared (IR) galaxies whose bolometric energy output is dominated in the IR ($\lambda$ = 8-1000 $\mu$m) by reprocessed dust emission (see \citealt{sanders96, chary01, franceschini01}). They were first catalogued in the local Universe ($z\leq$ 0.3) by the \textit{InfraRed Astronomical Satellite} (IRAS) and shown to only contribute $\sim$ 30\% to the CIRB which suggests that the majority originates from reprocessed dust emission by high redshift galaxies. The \textit{Spitzer Space Telescope} \citep{werner04} has greatly increased our understanding of the IR Universe at high redshifts. We now know that there is strong evolution in the population of IR galaxies with LIRGs dominating the cosmic star formation rate out to $z\sim1$ and ULIRGs out to $z\sim2$ \citep{lefloch05, perezgonzalez05, caputi07}. This evolution is also seen in the more luminous submillimeter galaxies (SMGs) at redshifts $>$ 2 \citep{smail97,scott02, clements08, dye08}. These studies have highlighted that a large fraction of star formation and gravitational accretion is heavily obscured by dust and therefore studying the properties and the nature of these galaxies is crucial to understanding galaxy evolution. \indent Deep surveys conducted by \textit{Spitzer}, particularly in the mid-IR (MIR) at 24 $\mu$m where the MIPS instrument is most sensitive have allowed us to understand the nature of dust obscured galaxies. Using spectroscopic diagnostics and/or modelling the spectral energy distributions (SEDs) combining optical and IR photometry has shown that star-formation is the dominant process with the AGN component becoming significant in the more luminous sources \citep{genzel00, farrah03, soifer08}. However the 24 $\mu$m surveys are limited as the peak of the SED for most IR luminous galaxies is in the range of $\sim 40 - 200$ $\mu$m. At high redshifts the 24 $\mu$m channel samples shorter wavelength regions contaminated by polycyclic aromatic hydrocarbon (PAH) emission and silicate absorption features making it difficult to obtain reliable estimates of the bolometric IR luminosities. \indent \textit{Spitzer} also observed at the longer 70 and 160 $\mu$m wavelengths and several studies have been carried out in order to characterise the 70 $\mu$m population. Analysis of the 160 $\mu$m population is severely hindered because of the lower resolution and sensitivity of the MIPS instrument at this wavelength. By modelling the IR SEDs of 70 $\mu$m selected galaxies in the $0.1 < z < 2$ redshift range, \cite{symeonidis07, symeonidis08} found that these objects require a cold (dust temperature $<$ 20K) emission component in the FIR. In the follow-up study, \cite{symeonidis09} (henceforth S09) suggest that the MIPS 70 $\mu$m population may be the missing link between the cold $z>1$ submillimeter (submm) population detected by SCUBA and the local IR luminous galaxies. \cite{kartaltepe10} extend their study to higher redshifts, $z\sim 3.5$ for 1503 reliable and unconfused 70 $\mu$m selected sources in the Cosmic Evolution Survey (COSMOS). Their analysis shows that the SED shapes are similar to local objects but also find evidence for a cooler component than is observed locally. Using ancillary radio and X-ray data, they find that the fraction of AGNs increases with L$_\mathrm{IR}$, and nearly 51\% of ULIRGs and all HLIRGs likely host a powerful AGN. \cite{symeonidis10} also arrive at the same conclusion although they find 23\% of ULIRGs contain an AGN. These figures are consistent with observations of local ULRIGs (see \citealt{veilleux95, sanders96, kim98}) which show that the AGN fraction increases from $\sim$ 4 \% at L$_\mathrm{IR}$ = 10$^{10}$ L$_\odot$ to $>$ 50\% at L$_\mathrm{IR} > $ 10$^{12}$ L$_\odot$. \indent In this paper, we examine the optical and IR properties from a spectroscopic follow-up of 293 70 $\mu$m selected sources from the XMM-LSS and Lockman Hole (LH) region of the \textit{Spitzer} Wide-area InfraRed Extragalactic survey (SWIRE) \citep{lonsdale03}. The present paper is the first in a series which aims to study the evolution of the 70 $\mu$m luminosity function. This paper is organised as follows: In Sect.\ref{sect:data} we present the sample selection criteria and in Sect.\ref{sect:obsdatareduc} we describe the observations and data reduction method. We present the main results in Sect.\ref{sect:results} and discuss the implications in Sect.\ref{sect:summary}. Throughout this paper, we assume a $\Lambda$CDM cosmology with $H_\mathrm{0}$ = 70kms$^{-1}$Mpc$^{-1}$, $\Omega_{\Lambda}$ = 0.7, and $\Omega_\mathrm{M}$ = 0.3. All magnitudes are in the AB system unless otherwise stated.

\label{sect:summary} We present spectroscopic follow-up observations of 70 $\mu$m selected sources in the SWIRE XMM-LSS and LH-ROSAT regions using the multi-object fibre spectrograph AF2/WYFFOS on WHT. We measured new spectroscopic redshifts for 293 70 $\mu$m and 35 24 $\mu$m sources. The redshift distribution for the 70 $\mu$m sources peaks at $z\sim0.3$ and has a high redshift tail out to $z\sim3.5$. The majority of the 24 $\mu$m selected QSOs are at typically $z>1$. The spectroscopic redshifts of our sample were compared with the photometric redshift estimates which show a good agreement for galaxies while some QSOs have poorer agreement which we attribute to optical variability or photometric redshift aliasing. \indent We carry out emission line diagnostic for 91 70 $\mu$m sources with [OIII], H$\beta$, [NII], H$\alpha$ and [SII] emission lines present to classify the galaxies into star-forming, Seyfert, composite and LINERs. We find in our sample 34 star-forming, 30 composite, 1 Seyfert, 5 LINERs and 21 ambiguous galaxies and 13 QSOs identified from their broad emission line spectra. \indent We then modelled the SEDs for each source from optical to the FIR using the method described in \cite{mrr05, mrr08} and SED templates from \cite{mrr08} . Our major finding is that the IR emission in the 70 $\mu$m sources are powered by star formation, with the AGN dust torus component making a small contribution for non QSOs. \cite{symeonidis10} examine the AGN content of a sample of IR luminous 70 $\mu$m selected galaxies by modelling the panchromatic SEDs from X-ray to IR and find for most of their sources, the AGN contributes less than 10\% to the IR budget. Moreover, they find all sources in their sample are primarily powered by star-formation. Our results are also consistent with the findings of \cite{trichas09} who find their X-ray selected sample with 70 $\mu$m detection are strongly star-forming and require a starburst SED to fit the 70 $\mu$m and 160 $\mu$m (if available) photometry. For 9 QSOs and the 1 Seyfert galaxy the AGN dust torus dominates in the NIR and MIR regimes producing from 40 to 100\% of the total IR luminosity. Of the 83 70 $\mu$m sources in our sample with 160 $\mu$m detections, 49 require the cirrus component which implies the presence of large amounts of cold dust at T$_d \sim$ 25 K. By integrating the SEDs from 8 - 1000 $\mu$m, we find our 70 $\mu$m sample is dominated by starbursts and LIRGs. We also find that the AGN fraction increases with L$_\mathrm{IR}$ with almost all (7/10) HLIRGs being optically identified as QSOs. \indent Finally, we fit single temperature modified blackbodies to 82 70 $\mu$m sources with 160 $\mu$m fluxes at $z<1.2$ with $\beta$ fixed at 1.5. Our sample has a mean and standard deviation dust temperatures, T$_{d}$ = 32.7 $\pm$ 6.7 K and dust mass range, M$_{d}$ = 10$^6$ - 10$^9$ M$_\odot$ and even increasing $\beta$ to 2 results in slightly lower temperatures, T$_{d}$ = 30 $\pm$ 6 K. In general we find that as $\beta$ increases T$_{d}$ decreases. Examining our sources in the L-T diagram shows that the 70 $\mu$m sources with 160 $\mu$m detection have dust temperatures that are systematically lower than the local IR luminous galaxies. \textit{Herschel} and \textit{Planck} will play a vital role in characterising the IR galaxy population, as well as breaking the degeneracy between $\beta$ and T$_{d}$ by providing data points across range of FIR/submm wavelengths.

1012.4168

1012

1012.1926_arXiv.txt

The large number of exoplanets found to orbit their host stars in very close orbits have significantly advanced our understanding of the planetary formation process. It is now widely accepted that such short-period planets cannot have formed {\em in situ}, but rather must have migrated to their current orbits from a formation location much farther from their host star. In the late stages of planetary formation, once the gas in the proto-planetary disk has dissipated and migration has halted, gas-giants orbiting in the inner disk regions will excite planetesimals and planetary embryos, resulting in an increased rate of orbital crossings and large impacts. We present the results of dynamical simulations for planetesimal evolution in this later stage of planet formation. We find that a mechanism is revealed by which the collision-merger of planetary embryos can kick terrestrial planets directly into orbits extremely close to their parent stars.

To date over 490 extrasolar planets have been discovered, revealing a wide diversity of planetary systems (http://exoplanet.eu). One of more unusual phenomena so revealed has been the population of ``Hot Jupiters'' -- gas-giants found in very small orbits (periods $<$ 8d) about their parent stars -- of which the prototype was the very first gas-giant exoplanet discovered, 51\,Peg \citep{may95}. It is believed that such short-period gas-giants cannot have formed this close to their parent stars, and so must have migrated in, or been scattered in, from a more distant formation region \citep{lin96,weid96,Ida04,cham09}. The measurement precisions that make the detection of such short-period exoplanets possible have over recent years continually improved for both Doppler (e.g. Gl 876 d \citep{rive05}, Gl 581 c \citep{udry07}, 61 Vir b \citep{vogt10}) and transit (e.g. Kepler-4b \citep{boru10}) detection. What then are the possible formation mechanisms that can produce such close-in terrestrial and super-terrestrial planets? Several models have been proposed for the formation of close-in terrestrial planets. \citet{raym06} have shown that super-Earths could form interior to a migrating Jovian planet. As they migrate inward, such gas-giants can shepherd planetary embryos interior to their orbits, which can then further collide and merge to generate Earth-like planets \citep{zhou05}. It has also been suggested that orbital migration and planet-planet scattering could potentially produce short-period super-Earths \citep{brun05, terq07, raym08}. Whatever the mechanism for their formation, it is likely that such planets are common around at least low-mass stars \citep{kenn08}. In all these scenarios, the formation of short-period Earth-like planets is associated with the migration of gas-giant planets. According to the core accretion paradigm for planetary formation, the isolation cores in the terrestrial planet formation region, and the solid cores of gas-giants, are both formed within $\sim 1$\,Myr from kilometer-sized planetesimals \citep{Saf69,wet80}. Subsequently massive solid cores accrete disk gas to form giant planets \citep{Kok02,Ida04} at $\sim 3-6$\,Myr, before the disk disperses \citep{hais01}. In the late stage of planet formation, when giant planets have ceased migration after the gas disk clears, the disk of countless planetesimals and planetary embryos will become turbulent due to stirring by gas-giants over hundreds of million years (or potentially even longer). In the meantime, it is expected that orbital crossings and giant impacts will frequently occur, which could lead to the formation of terrestrial planets \citep{cham01,raym04,zha09} and short-period Earth-like planets. In this Letter, we present a potential new formation mechanism for short-period Earth-like planets in the late stage of planet formation through a collision-merger scenario. In this mechanism, a planetary embryo is directly kicked to a close-in orbit after a collision with another embryo, and then the larger merged body is seized by the central star as a hot Earth-like planet.

We have uncovered a new mechanism for producing short-period terrestrial planets via collisions-mergers in the late stages of planetary formation. In this mechanism, two highly-eccentric bodies first undergo a severe orbital crossing and then form a short-period planet via collision-merger. In the set of simulations performed to date, this mechanism produces a short-period, terrestrial planet in 20$\%$ of runs. As mentioned previously, the formation rate for short-period terrestrial planets via a collision-merger process is only a moderate 20\%. However, this low rate may be a result of the limits imposed on our simulations by current computational capabilities, which restrict our adopted population of embryos and planetesimals to a few hundred objects with a total mass of only several times that of the Earth. The resultant planetesimal disk in our simulations is much smaller than that of the Minimum Mass Solar Nebula ($\sim0.01$ $M_{\odot}$ within 30\,AU \citep{weid77,haya81}) -- which would also contain billions of small bodies. Increasing the number of bodies and the total mass of the proto-planetary disk would likely increase the efficiency with which this mechanism produces short-period terrestrial planets. In addition, it is worth noting that close-in planets emerge from our simulations within a few million years. This is a significantly shorter timescale than the billion years over which the Solar System is thought to have undergone significant evolution. So, while near-infrared observations of young cluster samples, indicate an overall dust disk lifetime of $\sim6$\,Myr \citep{hais01}, the planetary system will actually continue to evolve over much longer timescales following the clearing of the dust and gas disk. During this late stage of planetary formation, frequent orbital crossings and huge impacts will occur, which are likely to significantly boost the feasibility of collision-merger events producing short-period terrestrial bodies. The collision-merger scenario for the formation of short-period planets does not require perfect accretion. Rather it relies on the collisions pushing the resultant body inward, so that the central star can capture it as a short-period planet. In this sense, such a mechanism could play a key role in throwing the largest fragments resulting from severe impacts into short-period orbits. On the other hand, given the diversity in the architectures of currently known systems, exoplanets are likely to form through a variety of mechanisms rather than through a uniform dominant process (D. Lin 2009, private communication). Our simulation results show one potential mechanism for the origin of short-period terrestrial planets in a compact disk with two gas-giants, and may predict an abundance of close-in bodies for this family.

1012.1926

1012

1012.3923_arXiv.txt

The growing number of extragalactic high-energy (HE, $E > 100$~MeV) and very-high-energy (VHE, $E > 100$~GeV) $\gamma$-ray sources that do not belong to the blazar class suggests that VHE $\gamma$-ray production may be a common property of most radio-loud Active Galactic Nuclei (AGN). In a previous paper, we have investigated the signatures of Compton-supported pair cascades initiated by VHE $\gamma$-ray absorption in monochromatic radiation fields, dominated by Ly$\alpha$ line emission from the Broad Line Region. In this paper, we investigate the interaction of nuclear VHE $\gamma$-rays with the thermal infrared radiation field from a circumnuclear dust torus. Our code follows the spatial development of the cascade in full 3-dimensional geometry. We provide a model fit to the broadband SED of the dust-rich, $\gamma$-ray loud radio galaxy Cen~A and show that typical blazar-like jet parameters may be used to model the broadband SED, if one allows for an additional cascade contribution to the {\it Fermi} $\gamma$-ray emission.

Blazars are a class of radio-loud active galactic nuclei (AGNs) comprised of Flat-Spectrum Radio Quasars (FSRQs) and BL~Lac objects. Their spectral energy distributions (SEDs) are characterized by non-thermal continuum spectra with a broad low-frequency component in the radio -- UV or X-ray frequency range and a high-frequency component from X-rays to $\gamma$-rays, and they often exhibit substantial variability across the electromagnetic spectrum. In the VHE $\gamma$-ray regime, the time scale of this variability has been observed to be as short as just a few minutes \citep{albert07,aharonian07}. While previous generations of ground-based Atmospheric Cherenkov Telescope (ACT) facilities detected almost exclusively high-frequency peaked BL~Lac objects (HBLs) as extragalactic sources of VHE $\gamma$-rays (with the notable exception of the radio galaxy M87), in recent years, a number of non-HBL blazars and even non-blazar radio-loud AGN have been detected by the current generation of ACTs. This suggests that most blazars might be intrinsically emitters of VHE $\gamma$-rays. According to AGN unification schemes \citep{up95}, radio galaxies are the mis-aligned parent population of blazars with the less powerful FR~I radio galaxies corresponding to BL~Lac objects and FR~II radio galaxies being the parent population of radio-loud quasars. Blazars are those objects which are viewed at a small angle with respect to the jet axis. If this unification scheme holds, then, by inference, also radio galaxies may be expected to be intrinsically emitters of VHE $\gamma$-rays within a narrow cone around the jet axis. While there is little evidence for dense radiation environments in the nuclear regions of BL~Lac objects --- in particular, HBLs ---, strong line emission in Flat Spectrum Radio Quasars (FSRQs) as well as the occasional detection of emission lines in the spectra of some BL~Lac objects \citep[e.g.,][]{vermeulen95} indicates dense nuclear radiation fields in those objects. This is supported by spectral modeling of the SEDs of blazars using leptonic models which prefer scenarios based on external radiation fields as sources for Compton scattering to produce the high-energy radiation in FSRQs, LBLs and also some IBLs \citep[e.g.,][]{ghisellini98,madejski99,bb00,acciari08}. If the VHE $\gamma$-ray emission is indeed produced in the high-radiation-density environment of the broad line region (BLR) and/or the dust torus of an AGN, it is expected to be strongly attenuated by $\gamma\gamma$ pair production \citep[e.g.][]{pb97,donea03,reimer07,liu08,sb08}. \cite{akc08} have suggested that such intrinsic $\gamma\gamma$ absorption may be responsible for producing the unexpectedly hard intrinsic (i.e., after correction for $\gamma\gamma$ absorption by the extragalactic background light) VHE $\gamma$-ray spectra of some blazars at relatively high redshift. A similar effect has been invoked by \cite{ps10} to explain the spectral breaks in the {\it Fermi} spectra of $\gamma$-ray blazars. This absorption process will lead to the development of Compton-supported pair cascades in the circumnuclear environment \citep[e.g.,][]{bk95,sb10,rb10}. In \cite{rb10}, we considered the full 3-dimensional development of a Compton-supported VHE $\gamma$-ray induced cascade in a monochromatic radiation field. This was considered as an approximation to BLR emission dominated by a single (e.g., Ly$\alpha$) emission line. In that work, we showed that for typical radio-loud AGN parameters rather small ($\sim \mu$G) magnetic fields in the central $\sim 1$~pc of the AGN may lead to efficient isotropization of the cascade emission in the {\it Fermi} energy range. We applied this idea to fit the {\it Fermi} $\gamma$-ray emission of the radio galaxy NGC~1275 \citep{abdo09a} under the plausible assumption that this radio galaxy would appear as a $\gamma$-ray bright blazar when viewed along the jet axis. In this paper, we present a generalization of the Monte-Carlo cascade code developed in \cite{rb10} to arbitrary radiation fields. In particular, we will focus on thermal blackbody radiation fields, representative of the emission from a circum-nuclear dust torus. In Section \ref{setup} we will outline the general model setup and assumptions and describe the modified Monte-Carlo code that treats the full three-dimensional cascade development. Numerical results for generic parameters will be presented in Section \ref{parameterstudy}. In Section \ref{CenA}, we will demonstrate that the broad-band SED of the radio galaxy Cen~A, including the recent {\it Fermi} $\gamma$-ray data \citep{abdo09c}, can be modeled with plausible parameters expected for a mis-aligned blazar, allowing for a contribution from VHE $\gamma$-ray induced cascades in the {\it Fermi} energy range. We summarize in Section \ref{summary}. \begin{figure}[ht] \centering \includegraphics[width=15cm]{f1.eps} \caption{\label{geometry}Geometry of the model setup.} \end{figure}

1012.3923

1012

1012.3324_arXiv.txt

The statistical analysis of the voids between galaxies is a field of research and the following catalogs have been explored: the two-degree Field Galaxy Redshift Survey (2dFGRS), see \cite{Patiri2006} and \cite{Benda-Beckmann2008}, the Sloan Digital Sky Survey (SDSS), see \cite{Tikhonov2007}, and the combination of 2dFGRS and SDSS, see \cite{Tinker2008}. The voids between galaxies are usually modelled by a survival function in the apparent radius as given by a modified exponential distribution, see (3) in \cite{Benda-Beckmann2008} or our Section~\ref{sectionself}; this fact is considered here a standard argument for comparison. The concept of Voronoi Diagrams dates back to the vortex theory applied to the solar system as developed in the 17th century, see \cite{Descartes1644} and Figure~1 in \cite{Aurenhammer2000}. The name is due to the two historical records by \cite{voronoi_1907,voronoi} and the applications to astrophysics beginning with ~\cite{kiang}. The Voronoi diagrams represent a model of the voids between galaxies, see \cite{Weygaert1989}, \cite{pierre1990}, \cite{barrow1990}, \cite{coles1991}, \cite{Weygaert1991a}, \cite{Weygaert1991b}, \cite{zaninettig}, \cite{Ikeuchi1991}, \cite{Subba1992}, \cite{Weygaert1994}, \cite{Goldwirth1995}, \cite{Weygaert2002}, \cite{Weygaert2003}, and \cite{Zaninetti2006}. The Poisson Voronoi tessellation (PVT) is a particular case of the Voronoi tessellation in which the seeds are generated independently on the $X$, $Y$ and $Z$ axes in 3D through a subroutine which returns a random real number taken from a uniform distribution between 0 and 1. For practical purposes, the subroutine RAN2 was used, see \cite {press}. On adopting an astrophysical point of view, the sectional PVT, $V_p(2,3)$, is very interesting because it allows a comparison with the voids as observed in slices of galaxies belonging to different catalogs such as the CFA2 catalog (\cite{geller}), the 2dfGRS (\cite{Norberg2002}), the 6dF Galaxy Survey (6dFGS) (\cite{Jones2004}) or the SDSS (\cite{Einasto2003}). The absence of (i) a numerical analysis through the survival function of normalized areas in 2D and normalized volumes in 3D of PVT (ii) a careful exploration of the statistical properties of $V_p(2,3)$, leads to the following questions being posed. \begin{itemize} \item Is it possible to integrate the usual probability density functions (PDFs) which characterize the main parameters of 2D and 3D PVT in order to obtain an analytical expression for the survival function? \item Is it possible to model the normalized areas of $V_p(2,3)$ with the known PDFs? \item Can we transform the normalized volumes and areas into equivalent radius distributions? \item Can we simulate the observed slices of galaxies as given, for example, by the 2dF Galaxy Redshift Survey? \end {itemize} In order to answer these questions, Section~\ref{secadopted} reports the three major PDFs used to model the normalized area/volume of 2D/3D PVT as well as the results of the fit. Section \ref{secradius} reports the apparent distribution in effective radius of the 3D PVT as well as their associated survival functions. Section \ref {sectionself} contains the self-similar survival function for voids between galaxies as well as the associated PDF. Section \ref{sectionsectional} reports the fit of the normalized area distribution of the sectional PVT with the Kiang function and the exponential distribution. Section \ref{seclarge} reports our actual knowledge of the photometric properties of galaxies as well as a Voronoi simulation. It is important to point out that the PVT is not used as a technique to identify voids in existing data catalogs, see \cite{Ebeling1993}, \cite{Bernardeau1996}, \cite{Schaap2000}, \cite{Marinoni2002}, \cite{Melnyk2006}, \cite{Schaap2009} and \cite{Elyiv2009}. The PVT formalism is here used conversely: to generate mock catalogs which are later calibrated by observations. On adopting the point of view of the statistical distributions is important to underline that the survival function is here identified with the cumulative void size distribution function. We briefly recall that the cumulative void size distribution function relates the number of voids to their size, analoguosly to the halo mass function which relates number of dark matter halos to their mass.

The PDFs which are usually used to model the normalized volume distribution of the 3D PVT are gamma type distributions such as the Kiang function, (\ref{kiang}), and the new Ferenc--Neda function, (\ref{rumeni}). Here, in order to make a comparison with the self-similar distribution of voids, we derived the survival distribution of the Kiang function, (\ref{survival_kiang}), and of the Ferenc--Neda function, (\ref{survivalfn3}). On adopting an astrophysical point of view the cut $V_p(2,3)$ may model the voids between galaxies as given by astronomical slices of the Millenium catalog. The analysis of the normalized areas of $V_p(2,3)$ is a subject of research rather than a well-established fact and we have fitted them with the Kiang function and the exponential distribution. The $\chi^2$ value indicates that the exponential distribution fits more closely the normalized area distribution of $V_p(2,3)$ than does the Kiang function, see Table~\ref{tablev23}. This fact follows from the comparison between the self-similar survival function and the exponential and Kiang distributions of the radius, see Figure~\ref{comparison_cut}. Therefore, the one parameter survival function of the radius of the exponential distribution for $V_p(2,3)$, $S_{ER23} $, as represented by (\ref{survival_expr23}), may model the voids between galaxies as well as the five parameter self-similar survival function. The observed 2dFGRS slices can be simulated but the behaviour of the luminosity function for galaxies and the consequent number of galaxies as a function of the redshift should be carefully analysed. We are planning, in future projects, \begin{itemize} \item To insert the thickness of the faces of PVT and to model the connected modification of the survival function. \item To apply the technique here developed to the real data sets, Millennium simulation, for example. The 3D nature of the method and detailed density mapping of the voids should be superior to other void identification algorithms as suggested by \cite{Schaap2009}. \end{itemize}

1012.3324

1012

1012.1321_arXiv.txt

The fringe sensor unit (FSU) is the central element of the phase referenced imaging and micro-arcsecond astrometry (PRIMA) dual-feed facility for the Very Large Telescope interferometer (VLTI). It has been installed at the Paranal observatory in August 2008 and is undergoing commissioning and preparation for science operation. Commissioning observations began shortly after installation and first results include the demonstration of spatially encoded fringe sensing and the increase in VLTI limiting magnitude for fringe tracking. However, difficulties have been encountered because the FSU does not incorporate real-time photometric correction and its fringe encoding depends on polarisation. These factors affect the control signals, especially their linearity, and can disturb the tracking control loop. To account for this, additional calibration and characterisation efforts are required. We outline the instrument concept and give an overview of the commissioning results obtained so far. We describe the effects of photometric variations and beam-train polarisation on the instrument operation and propose possible solutions. Finally, we update on the current status in view of the start of astrometric science operation with PRIMA.

\label{sec:intro} The PRIMA facility\cite{Delplancke:2008xr, van-Belle:2008eu} represents a major expansion of the VLTI infrastructure at Paranal Observatory in Chile. It introduces dual-feed capability and is designed for observation of two stellar objects within the isoplanatic angle of typically 0.5 arc-minutes. It consists of star-separator modules at the telescopes (STS\cite{Nijenhuis:2008cy}), differential delay lines (DDL\cite{Pepe2008}), an internal laser metrology\cite{Schuhler2007}, and the fringe sensor unit (FSU\cite{Sahlmann:2009kx}) in the beam-combination laboratory. The scientific drivers for the development of PRIMA are narrow-angle astrometry\cite{Launhardt2008} and phase referenced imaging in conjunction with the MIDI and AMBER instruments. The general status of the PRIMA facility is presented in Ref.~\citenum{Schmid2010}. Here, we present the status and results for the FSU. The FSU operates in the infrared $K$-band. It consists of two identical fringe sensors (FSUA and FSUB) and uses polarisation-dependent spatial encoding to obtain 4 simultaneous ABCD signals with relative phase-shifts of about 90 degrees. The beams are combined before spatial filtering by single-mode fibres to allow the injection of the internal laser metrology. A cryogenic low-resolution spectrograph disperses the light on 5 spectral pixels in each ABCD quadrant. The white-light pixel is not a physical one, but is synthesised from the sum of the 5 spectral pixels. Simultaneous estimates of group delay (GD) and phase delay (PD) are obtained from the spectral pixels and the white-light pixels, respectively. The FSU does not incorporate real-time photometric monitoring, i.e. there is no real-time access to the intensity in each interferometer arm or in the individual quadrants. For a detailed description of the instrument, the operation principle, and the observation setup, we refer the reader to Ref.~\citenum{Sahlmann:2009kx}. After concluding the laboratory tests in Europe\cite{Sahlmann2008a}, the FSU was installed at Paranal observatory in July and August 2008. First fringes and fringe tracking on a stellar object were obtained one month later. The FSU commissioning activities since then include the establishment of the observation range in terms of stellar magnitude and atmospheric conditions, the detailed assessment of the fringe tracking performance, and the automatisation of fringe detection and tracking. Caused by the un-availability of PRIMA subsystems, the commissioning of the PRIMA facility started later and first functional tests of the FSU together with other PRIMA subsystems were performed. Simultaneous fringe tracking on two stellar objects has not been achieved at the time of writing\cite{Schmid2010}, but is expected to be achieved within a few months. One important result obtained from on-sky fringe tracking is that the group delay and phase delay may not be robust against the perturbations that occur during observation. The FSU signals can be highly non-linear, which disturbs the fringe tracking loop and eventually affect the quality of planned scientific programs. Therefore, most of this contribution is concerned with the problem of non-linearity. \begin{figure}[t] \begin{center} \includegraphics[width= \linewidth, bb= -153 246 796 566]{12271f16} \end{center} \caption{Lock ratio over one minute (\textit{top row}) and residual OPD over one second (\textit{bottom row}) as function of \textit{K}-band magnitude, coherence time, seeing, and airmass. The Figure is from Ref.~\citenum{Sahlmann:2009kx}.} \label{fig:lr} \end{figure}

% The FSU has demonstrated spatially-encoded fringe tracking at the VLTI. The first results are encouraging, especially in terms of the limiting magnitude and tracking residuals, which is essential for the PRIMA scientific programmes. The remaining important parameter is the fringe tracking accuracy, which appears to be ambiguous during FSU fringe tracking. The problem is difficult to detect and quantify because a second, independent fringe sensor is required to unambiguously measure the tracking accuracy. Because the encountered problems can be caused by non-linear control signals, we have investigated the various sources of phase and group delay non-linearity. They are caused by errors that originate in the currently unavailable on-sky calibration of wavelengths and phase-shifts and in the inherent lack of real-time photometric monitoring capability of the FSU. The introduced perturbations equally affect the FSU estimates of visibility and signal-to-noise ratio, which is discussed in another contribution to these proceedings\cite{Schmid2010}. The assessment of the impact of non-linear control signals required the investigation of the fringe-tracking control system architecture. The effect of GD non-linearity on the fringe-tracking loop behaviour was examined using a simplified model. Some patterns appearing during on-sky fringe tracking could be reproduced by the model. Although this is not a proof that the model describes the real situation, it gives a possible explanation for the observed behaviour. We conclude by listing the necessary steps to characterise and eventually improve the system: \begin{itemize} \item Measure the FSU fringe-tracking accuracy with an independent fringe sensor attached to the system, which can be FINITO, the second FSU, MIDI, or AMBER. This will allow us to estimate the amplitude of the introduced biases. \item Measure the non-linearity of FSU phase and group delay during on-sky observation. This requires dual-feed observations (first dual-feed fringe tracking is planned for July 2010). \item Measure the phase-shifts and effective wavelengths during on-sky observation. In principle, this is possible by scanning the fringe packet fast enough to freeze the atmospheric turbulence, provided a precise delay measurement is available. Eventually, this will be part of the calibration plan for dual-feed astrometric observations and the procedure is described in\cite{Sahlmann:2009kx}. \item Based on the \emph{qualitative} assessment of non-linearity sources presented here, a \emph{quantitative} study has to be undertaken to identify the factors that are currently limiting the performance. \item A more sophisticated model of the control system that includes atmospheric noise, phase delay non-linearity and possibly atmospheric dispersion may be employed to study the fringe-tracking behaviour and its succeptibility to the various error sources in more detail. \item A better model for the group-delay measurement principle has to be developed, in order to study the impact of calibration errors and to optimise the algorithm's robustness. \item Implement upgraded algorithms for group and phase delay, that account for different ABCD wavelengths. \end{itemize} As the fringe sensor and the data producing entity, the FSU is the central element of the PRIMA facility. The comprehensive characterisation of the system are vital for the scientific exploration of the data. This process has to be finalised before the scientific programme of extrasolar planet search with PRIMA astrometry\cite{Launhardt2008} can begin.

1012.1321

1012

1012.2080.txt

Since photospheric bright points (BPs) were first observed, there has been a question as to how are they structured. Are they just single flux tubes or a bundle of the flux-tubes? Surface photometry of the quiet Sun (QS) has achieved resolution close to $0.1''$ with the New Solar Telescope at Big Bear Solar Observatory. This resolution allowed us to detect a richer spectrum of BPs in the QS. The smallest BPs we observed with TiO 705.68 nm were $0.13''$, and we were able to resolve individual components in some of the BPs clusters and ribbons observed in the QS, showing that they are composed of the individual BPs. Average size of observed BPs was $0.22''$.

Observations of the solar photosphere, so far, revealed a plethora of tiny bright features, usually concentrated in active regions or bordering the supergranules in the quiet Sun (QS). Near disk center, they appear as "Bright Points" (BPs), roundish or elongated bright features located in the intergranular dark lanes Dunn \& Zirker (1973), Mehltretter (1974), and Title {\it et al.} (1987).\par So far, the areas chosen for the study of BPs were active regions or plage because with prior resolution, the plethora of BPs was visible only there. The achieved resolution of the New Solar Telescope (NST) at Big Bear Solar Observatory (BBSO) revealed to us large number of BPs structures in the QS. The smallest observed BPs were around $0.1''$ leaving their substructure still a puzzling subject. Considering the recent result that BPs are possible source of the acoustic oscillations (Andic {\it et al.} 2010) the importance of their study increases.This work presents a details that might contribute to the answer of the possible substructure of the BP.

Starting with the ribbons that we detected in the quietest sun, we observed constant substructure. Even individual BPs that seem as a compact elements turned to be composed from one or more substructures that revealed themselves in course of the time. This happened for most of the BPs that were at least $0.5"$ in diameter. Our observation of BPs in the groups showed the tendency of the BPs to drift toward each other until they form a single structure. This result is similar to the observations made by Viticchie {\it et al.} (2009) in G-band line. Previously, it was reported that BPs could have various shapes (Dunn \& Zirker 1973, Mehltretter 1974, Title {\it et al.} 1987). However, most of the elongated shapes we observed were clusters of BPs, which during our time series unveiled their components. \par All BPs that had diameter close to the $0.13"$ showed low intensity when compared with the larger BPs. Since their size was matching our achieved diffraction limit we might speculate that in reality they are much smaller. And that their intensity is low because it fills our resolution element. \par To complement this result a signature of the flux-tubes collision inside a single bright point that did not separate itself has to be mentioned. Details concerning this finding are presented at the poster by A. Andic presented at this same conference. All this results encourage us to speculate that BPs, as we observe them now, are composed of the multiple flux tubes. P.R.Goode and NST are supported by following grands: NSF (AGS-0745744), NASA (NNX08BA22G) and AFOSR (FA9550-09-1-0655).

1012.2080

1012

1012.4243_arXiv.txt

{ We report the results of a successful 7~hour 1.4~GHz VLBI experiment using two new stations, ASKAP-29 located in Western Australia and WARK12M located on the North Island of New Zealand. This was the first geodetic VLBI observing session with the participation of these new stations. We have determined the positions of ASKAP-29 and WARK12M. Random errors on position estimates are 150--200~mm for the vertical component and 40--50~mm for the horizontal component. Systematic errors caused by the unmodeled ionosphere path delay may reach 1.3~m for the vertical component. }

\label{s:intro} As part of the Australian and New Zealand joint bid to host the multi-billion dollar Square Kilometre Array (SKA), both countries are investing heavily in advanced technologies for radio astronomy. In Australia this follows a strong tradition in radio astronomy and is expressed in the construction of the Australian SKA Pathfinder (ASKAP) on the western edge of the Australian continent \citep{r:jonston08}. Radio astronomy in New Zealand has links stretching back to the work of Elizabeth Alexander on solar radio emission \citep{r:ale46}. John Bolton and Gordon Stanley used a cliff interferometer to obtain rising and setting records of various radio sources \citep{r:bol82}; measurements made near Sydney and Auckland allowed them to identify mysterious ``radiostars'' with well known supernova remnants and galaxies \citep{r:bol49}. The first successful VLBI experiment between Australia and New Zealand was made in 2005 with a 6-m radio telescope near \Note{Auckland (Karaka) and} the Australia Telescope Compact Array (ATCA) \citep{r:gul05,r:tin06}. In 2008, Auckland University of Technology commissioned a new 12-m antenna at Warkworth near Auckland, --- the New Zealand's first research capable radio telescope \citep{r:gul09}. A joint Australia and New Zealand collaborative project is focused on developing both the ASKAP and Warkworth facilities as part of regional and global Very Long Baseline Interferometry (VLBI) arrays, for the purposes of both astronomy and geodesy. Due to the dearth of land mass in the Southern Hemisphere, both astronomical and geodetic VLBI has suffered in the past, with the radio telescopes used for VLBI restricted to the south east ``corner'' of the Australian continent, one telescope in South Africa and occasional use of radio telescopes in Antarctica and South America. \Note{The additional capability gained by adding radio telescopes located in Western Australia and New Zealand is substantial, increasing the angular resolution of the Australasian array by a factor of approximately four}, admitting a range of astronomy science goals described in \citet{r:jonston08}. High precision astrometry is likely to be an important part of the science case for the high angular resolution component of the SKA, in particular for precise determination of the distances of radio pulsars in the galaxy, to be used in various tests of fundamental physics \citep{r:smits10}. In order to achieve the astrometric performance required, the SKA will need to undertake astrometry utilising distributed clusters of small antennas operating as phased arrays. This is a departure from standard precision astrometry and our work using ASKAP as part of a VLBI array for astrometry will be an important testbed for the demonstration of these techniques. The 12-m radio telescope \wark\footnote{\wark\ and \askap\ are identifiers for specific VLBI antennas: near Warkworth, New Zealand and at the Murchison Radio-Astronomy Observatory, Western Australia, respectively.} is intended to be used as a part of the Australian Long Baseline Array (LBA), for spacecraft monitoring, and for VLBI observations in the framework of the International VLBI Service for Geodesy \& Astrometry (IVS)\footnote{{\tt http://ivscc.gsfc.nasa.gov/}} and the AuScope project\footnote{{\tt http://www.auscope.org.au/}}. \askap\ is an element of the ASKAP array of 36 identical dishes using Phased Array Feed (PAF) technologies, expected to be fully operational by 2013. ASKAP will undertake very wide field survey science in continuum and spectral line modes and is also intended to be used as a part of regional and global VLBI networks for a variety of projects\footnote{{\tt http://www.atnf.csiro.au/SKA/}}. As a part of commissioning the new antennas, the positions of antenna reference points should be determined. A reference point is defined as the point of the projection of the movable elevation axis onto the fixed azimuthal axis. For the analysis of VLBI source imaging experiments \Note{made in a phase-referencing mode}, the projection of a baseline vector, (i.e. vector between antenna reference points) to the tangential image plane should be known with errors not exceeding tens of centimetres, otherwise the image will be smeared \Note{\citet{r:cha2002}}. For astrometry applications, angular position accuracies of tens of $\mu$as are required, making the requirements on the accuracy of station positions much more stringent: 5--10~mm. One way of estimating the position of the antenna reference points is through analysis of a combination of a ground survey of markers attached to the antennas from a local network around the station and GPS observations from the points at the local network \citep[see e.g.,][]{r:sarti04,r:sarti09}. Another way to estimate station positions is to use the VLBI technique itself to determine group delays and then derive reference point positions from these group delay measurements. The advantage of this approach is that it also provides useful diagnostics on the VLBI equipment. The first fringe test experiment between \Askap\ and \wark\ was made in April 2010. First fringes on baselines between \askap\ and {\sc mopra} were found on April 22, 2010 and on baselines to \wark\ on the following day. This success prompted three first science experiments using the full Long Baseline Array (LBA) network with the participation of the two new stations: 1)~imaging observations of {\tt PKS 1934$-$638} on 29--30th April \citep{r:askap1}, 2)~a geodetic experiment on May 07, 2010, and 3)~imaging onbservations of {\tt Cen-A} on May 09, 2010 (Tingay et al., 2010, in preparation). We report here results from the geodetic experiment that was conducted at \askap\ and during the first geodetic observing session in May 07, 2010. The goal of this experiment was to determine the position of the antennas with decimeter accuracy and to collect diagnostic data. The characteristics of new antennas are presented in section \ref{s:ant}. The experiment and its analysis are described in sections \ref{s:exp} and \ref{s:anal}. Concluding remarks are made in section \ref{s:concl}.

\label{s:concl} We obtained the first estimates of the positions of \askap\ and \wark\ antenna reference points from VLBI observations. The random position $1\sigma$ errors, 5--6~cm for the horizontal coordinates and 20--40~cm for the vertical component, are close to that what one can expect from narrow-band observations at L~band. It is a pleasant surprise that the very first observations that followed the successful fringe test yielded a reasonable result. However, the systematic errors due to unaccounted ionosphere are significantly greater: 1.3~m for the vertical component and 0.14~m for the horizontal component. We also identified several problems with station equipment that will be fixed in the future. \Note{The use of a Rubidium frequency standard at \askap\ caused a decorrelation at a level of 4--5\% within 2 minute long scans. The frequency instability at longer time intervals appeared negligible with respect to errors caused by the ionosphere.} The existing L~band (1.4~GHz) receiver on \askap\ will be replaced \Note{in early 2011} with a new receiver comprising a purpose-built feed horn optimized for this antenna, with significantly higher efficiency and a noise calibration source. In 2011 \askap\ may be equipped with an X~band (8.4~GHz) receiver, and the possibility of an upgrade to a higher frequency receiver is being considered. \Note{In 2011, \wark\ was equipped with S/X dual-band receivers. Future geodetic experiments at these frequencies are planned for February 2011. The participation of stations \askap\ and \wark\ in VLBI experiments under absolute astrometry and geodesy programs will permit an improvement in the precision of results by two orders of magnitude and reach a millimeter level of accuracy. A more precise survey to provide 1--3 mm accuracy of the tie vector between VLBI and GPS antenna reference points is planned. Combined with further VLBI observations it will then be possible to reconcile the currently observed differences between GPS and VLBI positions of the \wark\ antenna reference point.} In November 2010, Telecom New Zealand made its 30-metre satellite Earth station available to AUT's Institute for Radio Astronomy and Space Research. It is a wheel-and-track beam-waveguide antenna built by NEC Corporation in 1984. It will be converted to a radio telescope capable of conducting both astronomical and geodetic research.

1012.4243

1012

1012.1466_arXiv.txt

We present ground-based and \textit{HST} optical observations of the optical transients (OTs) of long-duration Gamma Ray Bursts (GRBs) 060729 and 090618, both at a redshift of $z=0.54$. For GRB 060729, bumps are seen in the optical light curves (LCs), and the late-time broadband spectral energy distributions (SEDs) of the OT resemble those of local type Ic supernovae (SNe). For GRB 090618, the dense sampling of our optical observations has allowed us to detect well-defined bumps in the optical LCs, as well as a change in colour, that are indicative of light coming from a core-collapse SN. The accompanying SNe for both events are individually compared with SN1998bw, a known GRB-supernova, and SN1994I, a typical type Ic supernova without a known GRB counterpart, and in both cases the brightness and temporal evolution more closely resemble SN1998bw. We also exploit our extensive optical and radio data for GRB 090618, as well as the publicly-available \textit{SWIFT}-XRT data, and discuss the properties of the afterglow at early times. In the context of a simple jet-like model, the afterglow of GRB 090618 is best explained by the presence of a jet-break at $t-t_{o}>0.5$ days. We then compare the rest-frame, peak $V$-band absolute magnitudes of all of the GRB and X-Ray Flash (XRF)-associated SNe with a large sample of local type Ibc SNe, concluding that, when host extinction is considered, the peak magnitudes of the GRB/XRF-SNe cannot be distinguished from the peak magnitudes of non-GRB/XRF SNe.

Compelling evidence connecting the generation of bursts of gamma-rays and flashes of X-rays with the gravitational core-collapse of stripped-envelope supernovae (SNe) continues to grow, e.g., XRF 100316D and SN2010bh (Chornock et al. 2010; Starling et al. 2010). The connection between GRB 980425 and unusual type Ic SN1998bw (Galama et al. 1998) provided an intriguing clue to the so-called ``GRB-SN'' connection (e.g., Woosley \& Bloom 2006) that was later confirmed with the ``smoking-gun'' detection of GRB 030329 and SN2003dh (Stanek et al. 2003; Hjorth et al. 2003; Matheson et al. 2003). Supporting the spectroscopic detections of long-duration GRBs (i.e., burst durations greater than $\geq 2$ s) and XRF-associated SNe are numerous photometric detections, e.g., XRF 020903 (Bersier et al. 2006); GRB 041006 (Stanek et al. 2005). Late-time deviations from the power-law decline expected for a GRB afterglow, that are accompanied by a change in colour, are interpreted as evidence of supernovae (e.g., Zeh et al. 2004). This spectroscopic and photometric evidence appears to strongly favour massive-star models, such as the ``Collapsar'' model (Woosley 1993) and the millisecond-magnetar model (e.g., Usov 1992; Thompson 1994; Wheeler et al. 2000; Zhang \& M\' esz\' aros 2001), for GRB production. Indirect evidence supporting the ``GRB-SN'' connection was the realization that GRBs and type Ibc SNe occur in the brightest regions of their host galaxies (Fruchter et al. 2006; Kelly et al. 2008). These observations link GRBs and XRFs to sites of star-formation, and consequently to massive stars. Similar observations of the positions of Wolf-Rayet (WR) stars, a possible progenitor of GRBs and XRFs, also show that their distribution in their host galaxies favour the brightest regions (e.g., Leloudas et al. 2010). Additionally, observations of GRB host galaxy morphology reveal a high fraction of merging and interacting systems (e.g., Conselice et al. 2005; Wainwright et al. 2007), which was shown to be an efficient trigger of star formation in galaxies (e.g., Joseph et al. 1984; Kennicutt \& Keel 1984). The connection between long-duration GRBs and bright supernovae is not without complexity, however. While questions surround the nature of GRB 980425 in its relation to ``classical'' GRBs, two nearby events have challenged the idea that \textit{all} long-duration GRBs are accompanied by an optically-bright SN. GRBs 060505 and 060614 (Gal-Yam et al. 2006; Fynbo et al. 2006; Della Valle et al. 2006; Ofek et al. 2007) had no accompanying supernova down to limits more than 100 times fainter than SN1998bw (though the classification of GRB 060614 as a long-duration GRB is questionable (e.g., Gehrels et al. 2006; Norris et al. 2010). Thus the homogeneity of the ``GRB-SN'' connection, where the na\"ive expectation that \textit{all} long-duration GRB events produce enough nickel to power an optically-bright SN, may be an over-simplification of a much more complex reality. So while the ``GRB-SN'' connection has been established, many questions still remain such as: ``What kind of progenitors produce these events? Are the progenitors all the same? Why do some massive stars become GRB/XRFs while most do not?'' General understanding of the types of progenitors that can give rise to a GRB, as well as the environments they occur in, have provided indirect clues. The occurrence of ``optically dark bursts'' (i.e., GRBs without an optical afterglow) and the (sometimes) large amounts of dust in the vicinity of GRB progenitors that can obscure light from a supernova is a first attempt to explain the lack of SNe for some nearby GRBs, e.g., GRB 090417B (Holland et al. 2010). Additionally, nearby extremely faint (more than 100 times fainter than SN1998bw) core-collapse SNe have been detected (Zampieri et al. 2003; Valenti et al. 2009\footnote{Though see Foley et al. (2009) and Foley et al. (2010) who present various lines of evidence that SN 2008ha was the result of a thermonuclear explosion of a carbon-oxygen white dwarf.}) with low expansion velocities and extremely low nickel production. These under-luminous events provide another plausible explanation (e.g., Tominga et al. 2007; Moriya et al. 2010) for the lack of accompanying SNe with long-duration GRBs as well as illustrating the diverse optical properties of core-collapse SNe. The evidence gathered so far shows that at least \textit{some} long-duration GRBs and XRFs are accompanied by an optically-bright type Ic SN. Individual cases of spectroscopically-identified connections of type Ibc supernovae to GRBs and XRFs are the ideal way to further our understanding, however this is only technologically viable for low-redshift bursts. Circumstantial evidence provided photometrically has also furthered our understanding by highlighting the enigmatic differences in GRB/XRF-associated SNe. Here we present results of data obtained on \textit{HST} and ground-based facilities of GRB 060729, and ground-based data for GRB 090618, both of which are at a common redshift of $z = 0.54$. In section 2 we present our photometry and results for GRB 060729, effectively exploiting our \textit{HST} data to obtain image-subtracted magnitudes that provide evidence for an associated supernova. In section 3 we present our photometry and results for GRB 090618, which resulted in densely-sampled $R_{c}$ and $i$ band light curves (LCs) that clearly show ``bumps'' that are accompanied by a change in colour that we attribute to flux coming from a stripped-envelope, core-collapse supernova. In section 4 we compare our results of the rest-frame, absolute magnitudes of the SNe with those of: (1) all GRB/XRF-associated SNe and, (2) a sample of local type Ibc SNe, concluding that the progenitors that give rise to GRB/XRF SNe may be similar to those that produce local type Ibc SNe but with some differences. Throughout the paper, observer-frame times are used unless specified otherwise in the text. The respective decay and energy spectral indices $\alpha$ and $\beta$ are defined by $f_{\nu} \propto (t - t_{0})^{-\alpha}\nu^{-\beta}$, where $t_{0}$ is the time of burst and $\nu$ is the frequency. Foreground reddening has been corrected for using the dust maps of Schlegel et al. (1998), from which we find $E(B-V) =0.055$ mag for GRB 060729 and $E(B - V ) = 0.085$ mag for GRB 090618. We adopt a flat $\Lambda$CDM cosmology with a Hubble parameter of $H_{0} = 71$ km/s/Mpc, a matter density of $\Omega_{M} = 0.27$, and a cosmological constant of $\Omega_{\Lambda} = 1 - \Omega_{M} = 0.73$. For this cosmology a redshift of $z = 0.54$ corresponds to a luminosity distance of $d_{L} = 3098.7$ Mpc and a distance modulus of $42.45$ mag.

\subsection{The SNe accompanying GRBs 060729 \& 090618} We have presented photometric evidence for supernovae associated with GRBs 060729 and 090618. For GRB 060729, peak apparent magnitudes of the associated supernova are $R_{c} = 23.80 \pm 0.08$ and $I_{c} = 23.20 \pm 0.06$. When rest-frame extinction was considered using the analysis performed by Schady et al. 2010 ($E(B-V) \le 0.06$ mag), the peak, rest-frame absolute $V$-band magnitude was shown to be $M_{V} = -19.43 \pm 0.06$ which is approximately the same peak brightness as SN1998bw in the $V$-band. For GRB 090618 the peak apparent magnitudes of the associated supernova are $R_{c} = 23.45 \pm 0.08$ and $I_{c} = 23.00 \pm 0.06$. The rest-frame extinction local to the GRB was determined from the X-ray to optical SED fitting, and found to be $A_{V} = 0.3 \pm 0.1$. In turn these values imply a peak, rest-frame absolute $V$-band magnitude of the SN associated with GRB 090618 of $M_{V} = -19.75 \pm 0.13$, which is $\sim 0.3$ mag brighter than SN1998bw in the $V$-band. \subsection{Comparison with existing GRB \& XRF-associated SNe} \begin{figure} \centering \includegraphics[height=3.4in,width=2.5in, angle=270]{C_Frac.eps} \caption{Cumulative fraction plots for the two comparison scenarios: ($right$) between \textit{all} SNe and ($left$) considering only those events where an estimation of the host/rest-frame extinction has been made. In the latter case, the probabilities that the three sets of SNe are drawn from the same parent population are (1) GRB/XRF SNe \& All Ibc SNe: $P=0.88$, (2) GRB/XRF SNe \& only Ic SNe: $P=0.54$, which further supports the idea that GRB-SNe and type Ic SNe have similar progenitors.} \label{fig:CumulFrac} \end{figure} \begin{table*} \begin{minipage}{155mm} \caption{Peak Rest-Frame $V$-band Absolute Magnitudes for GRB \& XRF-producing SNe \label{table:GRB_SNe}} \begin{tabular}{ccccccc} \hline GRB & SN & Redshift ($z$) & $A_{V,foreground}$ & $A_{V,host}$ $^{a}$ & $M_{V}^{peak}$ (mag)$^{b,c}$ & Reference \\ \hline GRB 970228 & - & 0.695 & 0.543 & 0.15 & $-18.56 \pm 0.30$ & (1), (2), (3) \\ GRB 980326 & - & $\approx$ 1 & 0.26 & - & $\approx$ -19.5 & (4) \\ GRB 980425 & 1998bw & 0.0085 & 0.18 & 0.05 & $-19.42 \pm 0.30$ & (3), (5), (6), (7), (8), (31) \\ GRB 990712 & - & 0.434 & 0.09 & 1.67 & $-20.22 \pm 0.20$ & (3), (10), (11), (12), (31) \\ GRB 991208 & - & 0.706 & 0.05 & 0.76 & $-19.46 \pm 0.75$ & (9), (16) \\ GRB 000911 & - & 1.058 & 0.38 & 0.20 & $-18.31 \pm 0.15$ & (9), (16) \\ GRB 011121 & 2001ke & 0.36 & 1.33 & 0.39 & $-19.59 \pm 0.33$ & (3), (13), (14), (16) \\ GRB 020405 & - & 0.698 & 0.14 & 0.15 & $-19.46 \pm 0.25$ & (3), (15), (16), (31)\\ GRB 020410 & - & $\approx\ 0.5$ & 0.40 & 0.0 & $\approx -17.6$ & (3), (17) \\ XRF 020903 & - & 0.251 & 0.09 & 0.0 & $-18.89 \pm 0.30$ & (3), (18), (31)\\ GRB 021211 & 2002lt & 1.006 & 0.08 & 0.0 & $-18.27 \pm 0.60$ & (9), (19), (16)\\ GRB 030329 & 2003dh & 0.169 & 0.07 & 0.39 & $-19.14 \pm 0.25$ & (3), (16), (20), (31), (32) \\ XRF 030723 & - & $\approx\ 0.4$ & 0.089 & 0.23 & $\approx -17.9$ & (3), (9), (34) \\ GRB 031203 & 2003lw & 0.1055 & 2.77 & 0.85 & $-20.39 \pm 0.50$ & (3), (21), (22), (31) \\ GRB 040924 & - & 0.859 & 0.18 & 0.16 & $-17.47 \pm 0.48$ & (23) \\ GRB 041006 & - & 0.716 & 0.07 & 0.11 & $-19.57 \pm 0.30$ & (3), (16), (24)\\ GRB 050525A & 2005nc & 0.606 & 0.25 & 0.32 & $-18.76 \pm 0.28$ & (3), (25), (26), (33) \\ XRF 060218 & 2006aj & 0.033 & 0.39 & 0.13 & $-18.76 \pm 0.20$ & (3), (27), (28), (31)\\ GRB 060729 & - & 0.54 & 0.11 & 0.18 & $-19.43 \pm 0.06$ & This paper \\ GRB 080319B & - & 0.931 & 0.03 & 0.05 & $-19.12 \pm 0.40$ & (29), (33) \\ GRB 090618 & - & 0.54 & 0.27 & 0.3 & $-19.75 \pm 0.14$ & This paper \\ GRB 091127 & 2009nz & 0.49 & 0.12 & 0.0 & $-19.00 \pm 0.20$ & (30) \\ \hline \end{tabular} \medskip $^{a}$Host extinction where available. \\ $^{b}$Cosmological Parameters used: $H_{o} = 71$ km s$^{-1}$\ Mpc$^{-1}$, $\Omega_{M} = 0.27$, $\Omega_{\Lambda} = 0.73$. \\ $^{c}$Wherever errors are not quoted in the literature conservative errors of $0.4$ mag are used. \\ \ \\ (1) \cite{Galama00}, (2) \cite{CastLamb99}, (3) \cite{Richardson09}, (4) \cite{Bloom99}, (5) \cite{Galama98}, (6) \cite{McKSch99}, (7) \cite{Sollerman00}, (8) \cite{Nakamura01}, (9) \cite{Zeh04}, (10) \cite{Sahu00}, (11) \cite{Fruchter00}, (12) \cite{Christensen04a}, (13) \cite{Bloom02B}, (14) \cite{Garnavich03}, (15) \cite{Masetti03}, (16) \cite{Kann06}, (17) \cite{Levan05}, (18) \cite{Bersier06}, (19) \cite{DellaValle03}, (20) \cite{Matheson03}, (21) \cite{Malesani04}, (22) \cite{Mazzali06}, (23) \cite{Soderberg06}, (24) \cite{Stanek05}, (25) \cite{DellaValle06a}, (26) \cite{Blustin06}, (27) \cite{Sollerman06}, (28) \cite{Modjaz06}, (29) \cite{Tanvir10}, (30) \cite{Cobb10}, (31) \cite{Levesque10}, (32) \cite{Deng05}, (33) \cite{Kann10}, (34) \cite{Butler05}. \end{minipage} \end{table*} \begin{table*} \begin{minipage}{140mm} \caption{Peak Rest-Frame $V$-band Absolute Magnitudes for Local type Ibc \& Ic SNe \label{table:local_SNe}} \begin{tabular}{ccccccc} \hline Type & SN & Redshift ($z$) & $A_{V,foreground}$ & $A_{V,host}$ $^{a}$ & $M_{V}^{peak}$ (mag)$^{b,c}$ & Reference \\ \hline Ib & 1954A & 0.000977 & 0.07 & - & $-18.75 \pm 0.40$ & (1) \\ Ic & 1962L & 0.00403 & 0.12 & - & $-18.83 \pm 0.83$ & (2), (3)\\ Ic & 1964L & 0.002702 & 0.07 & - & $-18.38 \pm 0.65$ & (2), (4)\\ Ib & 1966J & 0.002214 & 0.04 & - & $-19.00 \pm 0.4$ & (4)\\ Ib & 1972R & 0.002121 & 0.05 & - & $-17.44 \pm 0.4$ & (5) \\ Ic & 1983I & 0.002354 & 0.04 & - & $-18.73 \pm 0.45$ & (2), (6)\\ Ib & 1983N & 0.001723 & 0.20 & 0.3 & $-18.58 \pm 0.57$ & (7)\\ Ib & 1983V & 0.005462 & 0.06 & 1.18 & $-19.12 \pm 0.41$ & (2), (8)\\ Ib & 1984I & 0.0107 & 0.33 & - & $-17.50 \pm 0.40$ & (9) \\ Ib & 1984L & 0.005281 & 0.08 & 0.0 & $-18.84 \pm 0.40$ & (10)\\ Ib & 1985F & 0.00167 & 0.06 & 0.63 & $-20.19 \pm 0.50$ & (11)\\ Ic & 1987M & 0.004419 & 0.08 & 1.28 & $-18.33 \pm 0.71$ & (2), (12), (13)\\ Ic & 1990B & 0.007518 & 0.10 & 2.53 & $-19.49 \pm 1.02$ & (2), (14)\\ Ib & 1991D & 0.041752 & 0.19 & 0.0 & $-20.01 \pm 0.60$ & (15)\\ Ic & 1991N & 0.003319 & 0.07 & - & $-18.67 \pm 1.06$ & (2)\\ Ic & 1992ar & 0.1451 & 0.30 & 0.0 & $-18.84 \pm 0.42$ & (2), (16)\\ Ic & 1994I & 0.001544 & 0.11 & 1.39 & $-17.49 \pm 0.58$ & (2), (17), (18)\\ Ic BL & 1997ef & 0.011693 & 0.13 & 0.55 & $-17.80 \pm 0.21$ & (2), (19), (34)\\ Ic pec & 1999as & 0.127 & 0.09 & 0.0 & $-21.21 \pm 0.20$ & (20)\\ Ib/c & 1999cq & 0.026309 & 0.16 & - & $-19.75 \pm 0.72$ & (2), (21)\\ Ib & 1999dn & 0.00938 & 0.16 & - & $-17.17 \pm 0.40$ & (22)\\ Ib/c & 1999ex & 0.011401 & 0.06 & - & $-17.67 \pm 0.26$ & (23)\\ Ib & 2001B & 0.005227 & 0.39 & - & $-17.13 \pm 0.40$ & (24) \\ Ic BL & 2002ap & 0.002187 & 0.29 & 0.0 & $-17.73 \pm 0.21$ & (2), (25)\\ Ic & 2003L & 0.021591 & 0.06 & - & $-18.90 \pm 0.40$ & (27)\\ Ic BL & 2003jd & 0.018826 & 0.14 & 0.29 & $-19.50 \pm 0.30$ & (19), (26), (34) \\ Ic & 2004aw & 0.0175 & 1.15 & 0.0 & $-18.05 \pm 0.39$ & (28)\\ Ic & 2004ib & 0.056 & 0.07 & - & $-16.94 \pm 0.40$ & (32) \\ Ib pec & 2005bf & 0.018913 & 0.14 & - & $-18.23 \pm 0.40$ & (33)\\ Ic BL & 2005fk & 0.2643 & 0.19 & - & $-20.41 \pm 0.40$ & (29)\\ Ic BL & 2005kr & 0.13 & 0.31 & 0.27 & $-19.08 \pm 0.40$ & (19), (29), (34)\\ Ic BL & 2005ks & 0.10 & 0.17 & 0.79 & $-18.41 \pm 0.40$ & (19), (29), (34)\\ Ib/c & 2007gr & 0.001728 & 0.19 & - & $-16.74 \pm 0.40$ & (30)\\ Ic BL & 2007ru & 0.01546 & 0.89 & 0.0 & $-19.09 \pm 0.20$ & (31)\\ \hline \end{tabular} \medskip $^{a}$Host extinction where available. \\ $^{b}$Cosmological Parameters used: $H_{o} = 71$ km s$^{-1}$\ Mpc$^{-1}$, $\Omega_{M} = 0.27$, $\Omega_{\Lambda} = 0.73$. \\ $^{c}$Wherever errors are not quoted in the literature conservative errors of $0.4$ mag are used. \\ \ \\ (1) \cite{Wild60}, (2) \cite{Richardson02}, (3) \cite{Bertola64}, (4) \cite{MillerBranch90}, (5) \cite{Barbon73}, (6) \cite{Tsvetkov83}, (7) \cite{Clocchiatti96}, (8) \cite{Clocchiatti97}, (9) \cite{Binggeli84}, (10) \cite{Tsvetkov87}, (11) \cite{Filippenko86}, (12) \cite{Filippenko90}, (13) \cite{Nomoto90}, (14) \cite{Clocchiatti01}, (15) \cite{Benetti02}, (16) \cite{Clocchiatti00}, (17) \cite{Yokoo94}, (18) \cite{Iwamoto94}, (19) \cite{Modjaz08}, (20) \cite{Hatano01}, (21) \cite{Matheson00}, (22) \cite{Qiu99}, (23) \cite{Martin99}, (24) BAOSS, (25) \cite{Mazzali02}, (26) \cite{Valenti08}, (27) \cite{Soderberg03}, (28) \cite{Taubenberger06}, (29) \cite{Baretine05}, (30) \cite{Foley07}, (31) \cite{Sahu09}, (32) \cite{Adelman05}, (33) \cite{Anupama05}, (34) \cite{Levesque10}. \end{minipage} \end{table*} \begin{table*} \begin{minipage}{165mm} \caption{Kolmogorov-Smirnov test results \label{Table:KS}} \begin{tabular}{ccccccc} \hline Dataset & Number of Data points & Mean & Standard Deviation & $P^{a}$ & $D^{a}$ & comments \\ \hline GRB/XRF-associated SNe & 22 & -19.02 & 0.77 & - & - & all events \\ Local type Ic SNe & 19 & -18.73 & 1.00 & 0.16 & 0.33 & all events \\ Local type Ibc SNe & 34 & -18.59 & 1.04 & 0.12 & 0.31 & all events \\ \hline GRB/XRF-associated SNe & 21 & -19.00 & 0.78 & - & - & only those with host $A_{V}$ \\ Local type Ic SNe & 12 & -18.75 & 1.03 & 0.54 & 0.27 & only those with host $A_{V}$ \\ Local type Ibc SNe & 17 & -18.93 & 0.97 & 0.88 & 0.18 & only those with host $A_{V}$ \\ \hline \end{tabular} \medskip $^{a}$Probability and maximum difference between the GRB/XRF SNe sample and the local SNe sample. \end{minipage} \end{table*} To put our detections into context, we compiled from the literature peak, rest-frame $V$-band absolute magnitudes for two samples of type Ibc SNe, incorporating the values of the host/rest-frame extinction (taken at the location of the GRB or SN when available; i.e., Kann et al. 2006; Modjaz et al. 2008; Levesque et al. 2010), and applying the SMC reddening law. The two samples are: (1) those associated with GRBs \& XRFs, and (2) local type Ibc SNe. We note that we have limited this analysis to incorporate only GRB/XRF events where an optically-bright SN has been positively detected. These samples, along with their respective references, are listed in Tables \ref{table:GRB_SNe} and \ref{table:local_SNe}. When comparing these two samples we are attempting to answer the question: ``Are the progenitors of GRB \& XRF associated SNe the same as those of local type Ibc SNe without an accompanying GRB/XRF trigger?'' by testing if the distribution of the peak magnitudes of the two samples of supernovae are different. To do this we performed a Kolmogorov-Smirnov (KS) test on the two samples. First we compared the entire GRB/XRF sample ($N=22$) with \textit{all} of the type Ibc SNe ($N=34$), finding a modest probability that the two samples are drawn from the same parent population of $P=0.12$. When we compare the GRB/XRF sample with only the type Ic SNe ($N=19$) we find a similar probability of $P=0.16$. However, when we limited the samples to include only those SNe where an estimation of the host/rest-frame extinction has been made, we see an increased probability between the datasets: $P=0.88$ between the GRB/XRF SNe ($N=21$) and all of the type Ibc SNe ($N=17$), and $P=0.54$ between the GRB/XRF SNe and only the type Ic SNe ($N=12$). For these samples a higher average peak magnitude is also seen among the local type Ibc SNe sample, with the local type Ibc SNe having $M_{V,Ibc} = -18.93$. This is an increase in average peak brightness of $\sim 0.3$ mag, and is likely due to the inclusion of five additional bright type Ib events with known host-extinction (SN: 1983N, 1983V, 1984L, 1985F \& 1991D) with the local Ic SNe sample (N=12). The average peak magnitude of these five SNe is $M_{V,Ib} = -19.34$, which places them among the brightest of the type Ib events. For comparison, a study by Richardson et al. (2002) analysed the absolute peak magnitudes of samples of all types of nearby SNe, and distinguished between ``Normal'' and ``Bright'' type Ibc events, with the latter having $M_{V,Ibc} = -19.72 \pm 0.24$, comparable with the average peak magnitudes of these five SNe. Thus the inclusion of these 5 bright Ib events has increased the average peak brightness of the type Ibc SNe sample, as well as increased the probability of association between our samples of host extinction-corrected type Ibc SNe and the GRB/XRF-SNe. The results of the K-S test are summarized in Table \ref{Table:KS} and cumulative fraction plots are shown in Figure \ref{fig:CumulFrac}. It is worth noting, however, that while the increase in probability may really be due to an increased association between the datasets, the caveat of smaller sample sizes is that as one goes to smaller samples it is harder to obtain a statistically significant discrepancy. Our analysis is not unique, with previous studies having performed similar analysis, which our results generally support. Richardson 2009 (R09) found from a sample of 14 GRB/XRF-SNe with (mostly) known values of the host extinction an average $M_{V,peak} = -19.2 \pm 0.2$, with $\sigma =0.7$, which our results are in agreement with. R09 also compared the GRB/XRF-SNe sample with a sample of stripped-envelope SNe, which included types Ib, Ic and IIb finding the GRB/XRF-SNe sample brighter by $\sim 0.8$ mag. We find that our complete GRB/XRF-SNe sample is $\sim 0.4$ mag brighter than the complete local type Ibc SNe (but only $\sim 0.1$ mag when only events with known host extinction are considered), which is somewhat less than R09, however our sample does not include any type IIb SNe. Ferrero et al. 2006 (F06) undertook a slightly different study, comparing the existing GRB/XRF-SNe sample (to date) with that of XRF 060218. They furthered an original analysis performed by Zeh et al. (2004), and calculated for each GRB-SN event the luminosity ratio $k$ and stretch factor $s$ in comparison with SN1998bw. They find for their host extinction-corrected sample of GRB/XRF-SNe clustering in the range $0.6 <\ k\ <\ 1.5$, implying a range of absolute magnitudes at peak brightness of $+0.55$ mag $ < \ M_{V}^{98bw}\ <\ -0.44$ mag, which is in good agreement with our results. They conclude that the width of the GRB/XRF-SN luminosity function is at least 2 mag, and there was no evidence that the luminosity function evolved with redshift. Our analysis concurs with both of these results. In conclusion, while the complete sample of GRB/XRF SNe are generally brighter ($M_{V,GRB} = -19.02$, $\sigma = 0.77$) than the complete local type Ibc SNe ($M_{V,Ibc} = -18.59$, $\sigma = 1.04$) and the type Ic SNe sample ($M_{V,Ic} = -18.73$, $\sigma = 1.00$), our test does not address factors such as host and progenitor metallicity and typical outflow velocities. Thus we cannot rule out the null-hypothesis that the samples of SNe are drawn from the same parent population. \ \\ \indent We are very grateful to Knut Olsen and Abi Saha for their assistance in obtaining images on the CTIO $4$-m telescope in August 2006. We would also like to thank the anonymous referee for their very thorough and constructive comments of the original manuscript. This work was supported partially by a Science and Technology Facilities Council (STFC) (UK) research studentship (ZC). MI, YJ, and WKP were supported by the Korea Science and Engineering Foundation (KOSEF) grant No. 2009-0063616, funded by the Korea government (MEST). The T40 telescope at the Observatorio de Aras de los Olmos is funded by the Generalitat Valenciana, the Spanish Ministerio de Ciencia e Innovacion and the Universitat de Valencia. The research of JG and AJCT is supported by the Spanish programmes ESP2005-07714-C03-03, AYA2007-63677, AYA2008-03467/ESP and AYA2009-14000-C03-01. TAF and ASM were supported by the grant of the President of the Russian Federation (MK-405.2010.2). The WSRT is operated by ASTRON with financial support from the Netherlands Organization for Scientific Research (NWO). AG acknowledges funding from the Slovenian Research Agency and from the Centre of Excellence for Space Sciences and Technologies SPACE-SI, an operation partly financed by the European Union, European Regional Development Fund and Republic of Slovenia, Ministry of Higher Education, Science and Technology. A.F.S. acknowledges support from the SpanishMICINN projects AYA2006-14056, Consolider-Ingenio 2007-32022, and from the Generalitat Valenciana project Prometeo 2008/132.

1012.1466

1012

1012.5715_arXiv.txt

We have performed the astrometry of H$_2$O masers in the Galactic star-forming region Onsala 2 North (ON2N) with the VLBI Exploration of Radio Astrometry. We obtained the trigonometric parallax of $0.261 \pm 0.009$ mas, corresponding to the heliocentric distance of $3.83 \pm 0.13$ kpc. ON2N is expected to be on the solar circle, because its radial velocity with respect to the Local Standard of Rest (LSR) is nearly zero. Using present parallax and proper motions of the masers, the Galactocentric distance of the Sun and the Galactic rotation velocity at the Sun are found to be $R_0 = 7.80\pm0.26$ kpc and $\Theta_0 = 213 \pm 5$ km s$^{-1}$, respectively. The ratio of Galactic constants, namely the angular rotation velocity of the LSR can be determined more precisely, and is found to be $\Omega_0=\Theta_0/R_0 = 27.3\pm0.8$ km s$^{-1}$ kpc$^{-1}$, which is consistent with the recent estimations but different from 25.9 km s$^{-1}$ kpc$^{-1}$ derived from the recommended values of $\Theta_0$ and $R_0$ by the International Astronomical Union (1985).

Very Long Baseline Interferometry (VLBI) astrometry is an important method to measure the structure of the Milky Way Galaxy (MWG). By measuring the accurate position of the source and its time variation, the source distance and proper motion can be determined directly. VLBI astrometry at 10 $\mu$as accuracy of the Galactic H$_2$O and CH$_3$OH maser sources with the VLBI Exploration Radio Astrometry (VERA) and the Very Long Baseline Array (VLBA) can determine accurate distances at kpc-scale with the errors less than 10\% (see for example \cite{hac06}; \cite{xu06}; \cite{hon07}). The galactocentric distance of the Sun, $R_0$, and the Galactic circular rotation velocity at the Sun, $\Theta_0$, are two fundamental parameters to study the structure of the MWG and they are called here as the Galactic constants. The rotation curve of the MWG and all kinematic distances of the sources in the MWG are derived from these parameters. Since 1985, the International Astronomical Union (IAU) has recommended to give the values of $R_0=8.5$ kpc and $\Theta_0=220$ km s$^{-1}$. However, recent studies report the values different from them (e.g. \cite{miyamoto98}; \cite{rei09a}). Estimation of the Galactic constants is, however, affected by several independent assumptions; the peculiar motion of the source, systematic non-circular motions of both the source and the Local Standard of Rest (LSR) due to the spiral arm potential, the non-axisymmetric potential of the MWG, the warping motion of the galactic disk, and, furthermore, the motion of the Sun with respect to the LSR. In this paper to simplify the situation we assume that the source moves with the perfect circular rotation on the disk. The circle with the raduis of the galactocentric distance $R_0$ of the Sun is called as the solar circle. All the sources on the circle perform the circular motion with the circular velocity $\Theta_0$ of the LSR, provided that the non-circular motion of a source in the galactic disk is negligible. Due to the symmetric geometry, the radial velocity of a source on the circle is observed to be zero with respect to the LSR, and the proper motion of the source depends only on the Galactic rotation velocity $\Theta_0$ of the LSR. Therefore, we can derive $\Theta_0$ from the measured proper motion of the source on the circle. We can also derive $R_0$ from the heliocentric distance of the source, since the source, the Sun, and the Galactic center make an isosceles triangle (Figure \ref{fig:4}). Thus, we can find directly the value of the angular velocity $\Omega_0$ of the LSR from $\Theta_0/R_0$. Traditionally, this value has been derived from the Oort constants $A$ and $B$ on the basis of the kinematic analysis of stars in the solar neighbourhood (see \cite{miyamoto98}). The value of this ratio is a constraint to estimate one of the Galactic constants from the other. Although the IAU gives recommended values of the Galactic constants, at least one of them should be revised, if the ratio is inconsistent to the observed value. Onsala 2 North (ON2N) is a massive star-forming region located at the Galactic coordinates of $(l, b) = (\timeform{75.78D}, \timeform{0.34D})$. Its radial velocity $v_{\mathrm{LSR}}$ with respect to the LSR is observed to be $0\pm1$ km s$^{-1}$ in the NH$_3$ and CS lines (\cite{olm99}; \cite{cod10}). \citet{lek06} detected the H$_2$O masers of ON2N at the radial velocity range from $-12$ to 9 km s$^{-1}$, with their peak flux densities of $10^2$--$10^3$ Jy. These H$_2$O masers are associated with a 7 mm radio continuum source and an NH$_3$ core as well, which are located at approximately 2" south from the ultracompact H\emissiontype{II} region, G75.78+0.34 (\cite{car97}; \cite{cod10}). Thus, ON2N is considered to be one of the sources on the solar circle. We made astrometric observations of the H$_2$O masers in ON2N with VERA. Based on these proper motions and trigonometric parallax measurement, we estimate the Galactic constants $\Theta_0$, $R_0$, and $\Omega_0$. In this paper we use the ``LSR velocity'' or $v_{\mathrm{LSR}}$ as the radial velocity with respect to the frame moving toward $\alpha_{1900}=\timeform{18h}$, $\delta_{1900}=\timeform{+30D}$ with $-20$ km s$^{-1}$ namely, a provisional solar motion with respect to the LSR to be determined after the traditional definition in radio astronomy since 1960s (see \cite{ker86}). \begin{figure} \begin{center} \FigureFile(80mm,80mm){./fig4.eps} \end{center} \caption{The method to estimate $R_0$ and $\Theta_0$ from the distance and the proper motion of the source on the solar circle.} \label{fig:4} \end{figure}

\subsection{Overall Properties of H$_2$O Masers in ON2N} Figure \ref{fig:1} shows the scalar-averaged cross-power spectra of H$_2$O masers in ON2N observed with the VERA Mizusawa-Iriki baseline on 2007/051 and 2008/104. The intense emissions with the flux density of $\geq 100$ Jy were detected at the LSR velocity range from $-5$ to 5 km s$^{-1}$. The center of the velocity range of the intense emissions is close to the LSR velocity of the associated molecular cloud at $v_{\rm LSR} = 0\pm1$ km s$^{-1}$ of the NH$_3$ and CS lines (\cite{olm99}; \cite{cod10}). The majority of $v_{\mathrm{LSR}}$ of H$_2$O masers is found in the range $-12 \leq v_{\rm LSR} \leq 9$ km s$^{-1}$ seen in the previous monitoring observations in 1995--2004 (\cite{lek06}). We also detected the blueshifted components at $v_{\rm LSR} = -33$ and $-17$ km s$^{-1}$ and the redshifted component at 19 km s$^{-1}$, which were not detected in the preveous observations (e.g. \cite{lek06}). Thirty H$_2$O maser spots were detected over the half of year and at more than three epochs. They were in the LSR velocity range from $-33$ to 19 km s$^{-1}$, and distrubuted with an area of 1.0"$\times$0.4". Figure \ref{fig:2} shows the distribution of internal motion of the maser spots in ON2N. The reference position of the map is set at the position of a maser spot at $v_{\rm LSR} = 0.1$ km s$^{-1}$, which is estimated to be $(\alpha, \delta)_{\rm J2000.0} = (\timeform{20h21m44.01225s}, \timeform{37D36'37.4844"})$. Although the H$_2$O masers are located at 2" south from the peak of the 6 cm radio continuum emission (\cite{woo89}), they are spatially coincident with the peaks of the 7 mm radio continuum emission and the NH$_3$ (3,3) emission (\cite{car97}; \cite{cod10}). \begin{figure*} \begin{center} \FigureFile(160mm,160mm){./fig1.eps} \end{center} \caption{A scalar-averaged cross-power spectra of H$_2$O masers in ON2N observed with the VERA Mizusawa-Iriki baseline at 2007/051 and 2008/104.} \label{fig:1} \end{figure*} \begin{figure*} \begin{center} \FigureFile(160mm,160mm){./fig2.eps} \end{center} \caption{(a): Distributions of H$_2$O masers (color filled circles) superimposed on the 6 cm radio continuum emission (cyan contour; \cite{woo89}), the 7 mm radio continuum emission (gray contour; \cite{car97}), and the NH$_3$ (3,3) emission (black contour; \cite{cod10}). (b): Internal motion vectors of H$_2$O masers. The spot color shows the LSR velocity. The arrow at the top left corner shows a internal motion of 1 mas yr$^{-1}$, corresponding to 18.2 km s$^{-1}$ at a distance of 3.83 kpc. (c): Close-up to the central part of (b).} \label{fig:2} \end{figure*} \subsection{Parallax and Proper Motion} The absolute motion of the respective H$_2$O maser spot, i.e. its motion with respective to the position reference source J2015+3710, is given by the sum of the proper motion and the annual parallactic motion. In order to separate these two motions we performed monitoring observations of the H$_2$O masers of ON2N for about two years. We made a combined parallax fit, which means a fitting of the positions of 30 H$_2$O maser spots to a common parallax, but different proper motion and position offset for each spot. Figure \ref{fig:3} and Table \ref{tab:1} show the results of the combined parallax fit. As can be clearly seen in Figure \ref{fig:3}, the observed points demonstrate a sinusoidal modulation with a period of 1 year caused by the annual parallax. For this fitting, we assigned independent ``error floors'' in quadrature with the formal position fitting uncertainties. Trial combined fits were conducted and a separate reduced $\chi^2$ (per degree of freedom) statistics was applied for the right ascension and declination residuals. The error floors of 0.088 mas and 0.111 mas in the right ascension and the declination respectively, were then adjusted iteratively so as to achive a reduced $\chi^2$ per dgree of freedom near unity in each coordinate. Combining all of the fittings, we obtained the trigonometric parallax of the H$_2$O maser spots $0.261\pm0.009$ mas. The parallax gives the heliocentric distance of ON2N $3.83\pm0.13$ kpc. In Table \ref{tab:1}, we also show the estimated parallaxes using individual fitting for each maser spot. We made this individual fitting only for 14 maser spots which were detected over one year. The obtained parallaxes from the individual fitting are consistent with each other and with the result of the combined fit. This means that the parallax obtained by the combined fitting is reliable. The systemic motion of the source can be estimated as an averaged motion of all maser spots, provided that internal motion is random or symmetric. We believe that this may be reasonable for ON2N because of the following two reasons. The averaged radial velocity of all maser spots is $-1.2$ km s$^{-1}$ which is close to the systemic radial velocity derived from the associated molecular cloud. This suggests that the maser spots move rather symmetrically. Figure \ref{fig:2} shows the residual proper motion vectors, which are differences between the individual proper motions and the average. We did not find any strong asymmetric motion. Therefore, we find that the reasonable absolute proper motions are not biased by the internal motions of ON2N. Thus, the systemic proper motion of ON2N is estimated to be $(\mu_{\alpha} \cos\delta, \mu_{\delta}) = (-2.79 \pm 0.13, -4.66 \pm 0.17)$ mas yr$^{-1}$ as the average motion of all 30 maser spots. Using the Galactic coordinates of ON2N $(l, b) = (\timeform{75.78D}, \timeform{-0.34D})$, the proper motion components in the galactic coordinates are calculated to be $(\mu_{l} \cos b, \mu_{b}) = (-5.42\pm0.16, -0.36 \pm 0.14)$ mas yr$^{-1}$ corresponding to the linear velocity of $(v_{l}, v_{b}) = (-98.4\pm2.9, -6.6 \pm 2.6)$ km s$^{-1}$ at the distance of 3.83 kpc. We note that these values are still affected by the solar motion, because the observed proper motion is not relative to the LSR but to the Sun. To convert this observed velocity to that with respect to the LSR, we have to fix the solar motion relative to the LSR. As mentioned in \S 1 we use the solar motion in the traditional definition of $(U_\odot, V_\odot, W_\odot)=(+10.3, +15.3, +7.7)$ km s$^{-1}$ after \citet{rei09a} \footnote {This velocity is consistent with the provisional solar motion we used, but the 3 dimensional linear velocity $(U_\odot, V_\odot, W_\odot)=(+10.0, +15.4, +7.8)$ km s$^{-1}$ shown in \citet{ker86} is inconsistent to the definition by themselves.}. We note that this set of values is very close to the one found by \citet{miyamoto98} (see Table 5) on the basis of the stellar motion analysis of the HIPPARCOS proper motions in the solar neighbourhood. Using this traditional solar motion, the observed proper motion is converted to $(\mu_{l} \cos b, \mu_{b}) = (-5.76 \pm0.16, 0.02 \pm0.14)$ mas yr$^{-1}$ and the corresponding linear velocity is $(v_{l}, v_{b}) = (-104.6 \pm2.9, 1.1 \pm 2.6)$ km s$^{-1}$ relative to the LSR. \begin{table*} \caption{The obtained values of parallax $\pi$ and proper motions $\mu_{\alpha \cos\delta}$ and $\mu_{\delta}$ for H$_2$O maser features in ON2N.} \label{tab:1} \begin{center} \begin{tabular}{rrrrcrrr} \hline & \multicolumn{1}{c}{$v_{\rm LSR}$} & \multicolumn{1}{c}{$\Delta \alpha \cos\delta$} & \multicolumn{1}{c}{$\Delta \delta$} & & \multicolumn{1}{c}{$\pi$} & \multicolumn{1}{c}{$\mu_{\alpha \cos \delta}$} & \multicolumn{1}{c}{$\mu_{\delta}$} \\ \multicolumn{1}{c}{ID} & \multicolumn{1}{c}{(km s$^{-1}$)} & \multicolumn{1}{c}{(mas)} & \multicolumn{1}{c}{(mas)} & \multicolumn{1}{c}{Epochs} & \multicolumn{1}{c}{(mas)} & \multicolumn{1}{c}{(mas yr$^{-1}$)} & \multicolumn{1}{c}{(mas yr$^{-1}$)} \\ \hline 1& $-32.8$ &$ 3.0$&$ 46.7$&\texttt{........IJK}&----- &$-1.53\pm0.24$&$-2.12\pm0.31$\\ 2& $-17.6$ &$ -39.5$&$ -28.1$&\texttt{ABCDEFG....}&----- &$-2.89\pm0.14$&$-5.04\pm0.17$\\ 3& $-16.5$ &$ -39.0$&$ -28.2$&\texttt{ABCDEFGH...}&0.252$\pm$0.043&$-3.67\pm0.10$&$-5.09\pm0.13$\\ 4& $-12.1$ &$ -15.3$&$ -16.7$&\texttt{ABCDEF.....}&----- &$-3.12\pm0.21$&$-4.82\pm0.26$\\ 5& $-10.4$ &$-165.9$&$ -2.4$&\texttt{....EFGHIJ.}&0.258$\pm$0.024&$-4.96\pm0.10$&$-4.32\pm0.13$\\ 6& $ -4.8$ &$ -0.2$&$ -2.6$&\texttt{......G.IJK}&----- &$-2.76\pm0.12$&$-5.16\pm0.16$\\ 7& $ -4.8$ &$ 4.8$&$ 5.5$&\texttt{ABCDEFGH...}&0.236$\pm$0.036&$-3.01\pm0.10$&$-4.97\pm0.13$\\ 8& $ -3.7$ &$ 4.6$&$ 5.4$&\texttt{ABCDEFGHI..}&0.257$\pm$0.037&$-2.75\pm0.08$&$-5.08\pm0.10$\\ 9& $ -1.8$ &$-146.6$&$ 0.6$&\texttt{ABCDEFGHIJK}&0.229$\pm$0.038&$-4.90\pm0.04$&$-4.15\pm0.06$\\ 10& $ -1.2$ &$ 0.2$&$ 0.0$&\texttt{ABCDEFGHIJK}&0.276$\pm$0.031&$-2.67\pm0.05$&$-4.58\pm0.06$\\ 11& $ 0.1$ &$ 0.0$&$ 0.0$&\texttt{ABCDEFGHIJK}&0.287$\pm$0.039&$-2.59\pm0.05$&$-4.73\pm0.06$\\ 12& $ 0.5$ &$ 626.1$&$-185.4$&\texttt{.......HIJK}&----- &$-1.62\pm0.15$&$-4.77\pm0.19$\\ 13& $ 0.7$ &$ 625.2$&$-187.6$&\texttt{.......HIJK}&----- &$-1.88\pm0.15$&$-4.84\pm0.19$\\ 14& $ 0.7$ &$ 626.2$&$-180.2$&\texttt{ABCDEFG....}&----- &$-1.50\pm0.14$&$-5.04\pm0.17$\\ 15& $ 1.4$ &$-439.8$&$ 216.2$&\texttt{ABCDEFGHIJK}&0.252$\pm$0.031&$-3.51\pm0.05$&$-4.05\pm0.06$\\ 16& $ 1.4$ &$ 623.2$&$-187.9$&\texttt{........IJK}&----- &$-1.62\pm0.24$&$-4.68\pm0.30$\\ 17& $ 1.6$ &$ 627.7$&$-181.7$&\texttt{ABCDEFGHIJK}&0.272$\pm$0.048&$-2.17\pm0.05$&$-4.62\pm0.06$\\ 18& $ 1.8$ &$-439.5$&$ 216.4$&\texttt{ABCDEFGHIJK}&0.275$\pm$0.031&$-3.67\pm0.05$&$-4.25\pm0.06$\\ 19& $ 1.8$ &$ 625.1$&$-187.3$&\texttt{........IJK}&----- &$-1.49\pm0.24$&$-4.73\pm0.31$\\ 20& $ 1.8$ &$ 625.8$&$-181.2$&\texttt{ABCDEFGHIJK}&0.250$\pm$0.034&$-2.03\pm0.05$&$-5.23\pm0.06$\\ 21& $ 2.0$ &$-450.5$&$ 239.4$&\texttt{ABCDEFG.IJK}&0.264$\pm$0.038&$-3.07\pm0.05$&$-4.27\pm0.06$\\ 22& $ 2.0$ &$-439.1$&$ 218.2$&\texttt{ABCDEFGH...}&0.270$\pm$0.043&$-3.85\pm0.10$&$-3.56\pm0.13$\\ 23& $ 2.2$ &$ -17.1$&$ -18.7$&\texttt{..CDEF.....}&----- &$-3.66\pm0.35$&$-5.07\pm0.44$\\ 24& $ 2.4$ &$ -15.8$&$ -17.4$&\texttt{.....FGHIJK}&----- &$-2.37\pm0.09$&$-4.86\pm0.11$\\ 25& $ 2.6$ &$ -17.1$&$ -18.4$&\texttt{......GHIJK}&----- &$-2.39\pm0.11$&$-4.87\pm0.15$\\ 26& $ 3.3$ &$-409.6$&$ 160.2$&\texttt{.......HIJK}&----- &$-3.25\pm0.15$&$-4.18\pm0.19$\\ 27& $ 5.8$ &$ -9.7$&$ -5.1$&\texttt{.BCDEFGHI..}&0.292$\pm$0.050&$-2.96\pm0.09$&$-4.31\pm0.11$\\ 28& $ 8.3$ &$ -1.9$&$ -1.2$&\texttt{ABCDEF.....}&----- &$-2.61\pm0.20$&$-4.44\pm0.26$\\ 29& $ 8.7$ &$ -52.7$&$ -7.2$&\texttt{ABCDEF.....}&----- &$-2.84\pm0.21$&$-5.10\pm0.25$\\ 30& $ 19.3$ &$ 39.5$&$ 10.6$&\texttt{.BCDEF.....}&----- &$-2.36\pm0.27$&$-6.91\pm0.34$\\ \hline \multicolumn{5}{c}{Combined fit} & $0.261 \pm 0.009$ & & \\ \multicolumn{5}{c}{Average} & & $-2.79 \pm 0.13$ & $-4.66 \pm 0.17$ \\ \hline \multicolumn{8}{@{}l@{}} {\hbox to 0pt{\parbox{150mm}{\footnotesize Column (3), (4): Right ascension and declination offsets relative to the positon of the maser spot at $v_{\rm LSR} = 0.1$ km s$^{-1}$, and $(\alpha, \delta)_{\rm J2000.0} = (\timeform{20h21m44.01225s}, \timeform{37D36'37.4844"})$.\\ Column (5): Each alphabetical letter represents the epoch with maser detection. A, B, C, ..., and K mean the 11 epochs from 53847, 53941, ..., and 54657 in MJD, respectively. A dot represents the epoch without the detection. \\ Column (6): Parallax estimated from the individual fitting.\\ Column (7), (8): Motions on the sky in the directions along the right ascension and declination.\\ }\hss}} \end{tabular} \end{center} \end{table*} \begin{figure*} \begin{center} \FigureFile(160mm,80mm){./fig3.eps} \end{center} \caption{Parallax of the H$_2$O masers in ON2N. The data for the different maser positions are slightly shifted in time for clarity. Individual proper motions and position offsets are removed. The left and right panels show the annual parallax in right ascension and declination, respectively. The numbers in each panel show the ID number of the spot listed in Table \ref{tab:1}.} \label{fig:3} \end{figure*} \subsection{Derivation of the Galactic Constants} Based on the observed LSR velocity of ON2N we believe that the source is located at or close to the solar circle. The small proper motion along the galactic latitude, $v_b=1.1\pm2.6$ km s$^{-1}$, shown in \S 3 supports that ON2N rotates circularly around the Galactic center. For an object on the solar circle, the galactocentric distance of the Sun or that of the source, $R_0$, is estimated from the heliocentric distance of the source as \begin{equation} R_0 = \frac{D}{2 \cos l}, \label{equ:1} \end{equation} where $l$ is the galactic longitude of the source (see Figure \ref{fig:4}). Our estimation of the heliocentric distance of ON2N yields $D=3.83\pm0.13$ kpc. It gives $R_0=7.80\pm0.26$ kpc, if ON2N is exactly located on the solar circle. This value is close to the previous estimations; $R_0$ is estimated to be $8.0\pm0.5$ kpc from a combination of many methods reviewd by \citet{rei93}, $7.9^{+0.8}_{-0.7}$ kpc from a parallax measurement of H$_2$O masers in Sgr B2 with VLBA (\cite{rei09b}), and $8.28\pm0.44$ kpc using the orbits of stars around Sgr A* from VLT and Keck data (\cite{gil09}). For a source on the solar circle with purely circular rotation, its proper motion velocity along the galactic longitude, $v_l$, gives the galactic rotation velocity of the LSR or that of the object as \begin{equation} \Theta_0 = - \frac{v_l}{2 \cos l}, \label{equ:2} \end{equation} where $l$ is the galactic longitude of the source. Our measured value of $v_l=-104.6\pm2.6$ km s$^{-1}$ gives $\Theta_0=213\pm5$ km s$^{-1}$. This value is smaller than that estimated by \citet{rei09a}, $\Theta_0=254\pm16$ km s$^{-1}$ but close to the IAU recommended value of $\Theta_0=220$ km s$^{-1}$. Our derivations of the Galactic constants are strongly dependent on the assumption of the location of ON2N in the MWG. However, we found that the ratio of the Galactic constants, $\Theta_0 / R_0$, which is the anglular velocity of the LSR, $\Omega_0$, can be estimated less dependently on the assumption. The ratio $\Theta_0 / R_0$ can be estimated, even if ON2N is not exactly located on the solar circle, but near there. For a source on pure circular rotation at any position in the galactic disk, its radial and tangential velocities with respect to the LSR can be written as \begin{eqnarray} v_r &=& \left( \frac{\Theta}{R} - \frac{\Theta_0}{R_0} \right) R_0 \sin l, \label{equ:3} \\ v_l &=& \left( \frac{\Theta}{R} - \frac{\Theta_0}{R_0} \right) R_0 \cos l - \frac{\Theta}{R}D, \label{equ:4} \end{eqnarray} where $R$ is the actual galactocentric distance of the source, and $D$ is the heliocentric distance of the source, and $\Theta$ is the galactic rotation velocity of the source. Equations (\ref{equ:3}) and (\ref{equ:4}) yield \begin{eqnarray} \nonumber \frac{\Theta_0}{R_0} &=& -\frac{v_l}{D} + v_r \left(\frac{1}{D \tan l} - \frac{1}{R_0 \sin l}\right) \\ &=& -a_0 \mu_l + v_r \left(\frac{1}{D \tan l} - \frac{1}{R_0 \sin l}\right), \label{equ:7} \end{eqnarray} where $a_0$ is a constant to convert the unit from an angular velocity to a linear velocity and $a_0=4.74$ km s$^{-1}$ mas$^{-1}$ yr kpc$^{-1}$. For a source near the solar circle, its $v_r$ is nearly zero. In this case, equation (\ref{equ:7}) yields \begin{equation} \frac{\Theta_0}{R_0} \simeq -a_0 \mu_l, \end{equation} which is free from $R_0$. Actually we found that $\Theta_0$/$R_0$ is a nearly constant at $6 \leq R_0 \leq 10$ kpc and the ratio is obtained as $\Theta_0 / R_0 = 27.3 \pm 0.8$ km s$^{-1}$ kpc$^{-1}$ using the $D = 3.83\pm0.13$ kpc, $\mu_l = -5.76 \pm 0.16$ mas yr$^{-1}$, and $v_r = 0 \pm 1$ km s$^{-1}$. This value is close to the value of $\Theta_0 / R_0 = 28.7 \pm 1.3$ km s$^{-1}$ kpc$^{-1}$ obtained from a tangent point source, ON1 (\cite{nag11}), and the value of $\Theta_0 / R_0 = 28.6 \pm 0.2$ km s$^{-1}$ kpc$^{-1}$ obtained from the proper motion measurement of Sgr A* (\cite{rei04}), which is revised using the traditional definition of the solar motion by us. However, this value is inconsistent to that derived from the IAU recommended values 220 km s$^{-1}$ / 8.5 kpc=25.9 km s$^{-1}$ kpc$^{-1}$. This estimation gives another constraint on the Galactic constants which is independent of the Oort constants derived from stellar motion near the Sun. \bigskip We thank the referee Dr. Masanori Miyamoto for his invaluable comments and suggestions. We also thank to the staff members of all the VERA stations for their assistances in the observations.

1012.5715

1012

1012.5186_arXiv.txt

We present the first direct and unbiased measurement of the evolution of the dust mass function of galaxies over the past 5 billion years of cosmic history using data from the Science Demonstration Phase of the {\em Herschel}-ATLAS. The sample consists of galaxies selected at 250\mic which have reliable counterparts from SDSS at $z<0.5$, and contains 1867 sources. Dust masses are calculated using both a single temperature grey-body model for the spectral energy distribution and also using a model with multiple temperature components. The dust temperature for either model shows no trend with redshift. Splitting the sample into bins of redshift reveals a strong evolution in the dust properties of the most massive galaxies. At $z=0.4-0.5$, massive galaxies had dust masses about five times larger than in the local Universe. At the same time, the dust-to-stellar mass ratio was about 3--4 times larger, and the optical depth derived from fitting the UV--sub-mm data with an energy balance model was also higher. This increase in the dust content of massive galaxies at high redshift is difficult to explain using standard dust evolution models and requires a rapid gas consumption timescale together with either a more top-heavy IMF, efficient mantle growth, less dust destruction or combinations of all three. This evolution in dust mass is likely to be associated with a change in overall ISM mass, and points to an enhanced supply of fuel for star formation at earlier cosmic epochs.

The evolution of the dust content of galaxies is an important and poorly understood topic. Dust is responsible for obscuring the UV and optical light from galaxies and thus introduces biases into our measures of galaxy properties based on their stellar light (Driver et al. 2007). The energy absorbed by dust is re-emitted at longer infra-red and sub-millimetre (sub-mm) wavelengths, providing a means of recovering the stolen starlight. Dust emission is often used as an indicator of the current star formation rate in galaxies - although this calibration makes the assumption that young, massive stars are the main source of heating for the dust and that the majority of the UV photons from the young stars are absorbed and re-radiated by dust (Kennicutt et al. 1998, 2009; Calzetti et al. 2007). Many surveys of dust emission from 24--850\mic\ (Saunders et al. 1990; Blain et al. 1999; Le Floc'h et al. 2005; Gruppioni et al. 2010; Dye et al. 2010; Eales et al. 2010) have noted the very strong evolution present in these bands and this is usually ascribed to a decrease in the star formation rate density over the past 8 billion years of cosmic history ($z\sim1$: Madau et al. 1996, Hopkins 2004). The interpretation of this evolution is complicated by the fact that the dust luminosity of a galaxy is a function of both the dust content and the temperature of the dust. It is pertinent to now ask the question ``{\em What drives the evolution in the FIR luminosity density?\/}'', is it an increase in dust heating (due to enhanced star formation activity) or an increase the dust content of galaxies (due to their higher gas content in the past) -- or both? Dust is thought to be produced by both low-intermediate mass AGB stars (Gehrz 1989; Ferrarotti \& Gail 2006; Sargent et al. 2010) and by massive stars when they explode as supernovae at the end of their short lives (Rho et al. 2008; Dunne et al. 2009; Barlow et al. 2010). Thus, the dust mass in a galaxy should be related to its current and past star formation history. Dust is also destroyed through astration and via supernovae shocks (Jones et al. 1994), and may also reform through accretion in both the dense and diffuse ISM (Zhukovska et al. 2008; Inoue 2003; Tielens 1998). The life cycle of dust is thus a complicated process which many have attempted to model (Morgan \& Edmunds 2003; Dwek et al. 1998; Calura et al. 2008, Gomez et al. 2010; Gall, Anderson \& Hjorth 2011) and yet the basic statistic describing the dust content of galaxies - the dust mass function (DMF) - is not well determined. The first attempts to measure the dust mass function were made by Dunne et al. (2000; hereafter D00) and Dunne \& Eales (2001; hereafter DE01) as part of the SLUGS survey using a sample of {\em IRAS} bright galaxies observed with SCUBA at 450 and 850\mic. Vlahakis, Dunne \& Eales (2005; hereafter VDE05) improved on this by adding an optically selected sample with sub-mm observations. These combined studies, however, comprised less than 200 objects - none of which were selected on the basis of their dust mass. These studies were also at very low-z and did not allow for a determination of evolution. A high-z dust mass function was estimated by Dunne, Eales \& Edmunds (2003; hereafter DEE03) using data from deep sub-mm surveys. This showed considerable evolution with galaxies at the high mass end requiring an order of magnitude more dust at $z\sim 2.5$ compared to today (for pure luminosity evolution), though with generous caveats due to the difficulties in making this measurement. Finally, Eales et al. (2009) used BLAST data from 250--500\mic\ and also concluded that there was strong evolution in the dust mass function between $z=0-1$ but were also limited by small number statistics and confusion in the BLAST data due to their large beam size. In this paper, we present the first direct measurement of the space density of galaxies as a function of dust mass out to $z=0.5$. Our sample is an order of magnitude larger than previous studies, and is the first which is near `dust mass' selected. We then use this sample to study the evolution of dust mass in galaxies over the past $\sim 5$ billion years of cosmic history in conjunction with the elementary dust evolution model of Edmunds (2001). The new sample which allows us to study the dust mass function in this way comes from the {\em Herschel}-Astrophysical Terahertz Large Area Survey (H-ATLAS; Eales et al., 2010), which is the largest open-time key project currently being carried out with the {\em Herschel} Space Observatory (Pilbratt et al., 2010). H-ATLAS will survey in excess of 550 deg$^2$ in five bands centered on 100, 160, 250, 350 and 500$\mu$m, using the PACS (Poglitsch et al., 2010) and SPIRE instruments (Griffin et al., 2010). The observations consist of two scans in parallel mode reaching 5$\sigma$ point source sensitivities of 132, 126, 32, 36 and 45 mJy in the 100, 160, 250, 350 \&\ 500$\mu$m bands respectively, with beam sizes of approximately 9\asec, 13\asec, 18\asec, 25\asec\ and 35\asec. The SPIRE and PACS map-making are described in the papers by Pascale et al. (2011) and Ibar et al. (2010), while the catalogues are described in Rigby et al. (2011). One of the primary aims of the {\em Herschel}-ATLAS is to obtain the first unbiased survey of the local Universe at sub-mm wavelengths, and as a result was designed to overlap with existing large optical and infrared surveys. These Science Demonstration Phase (SDP) observations are centered on the 9$^h$ field of the Galaxy And Mass Assembly (GAMA; Driver et al. 2011) survey. The SDP field covers 14.4 sq. deg and comprises approximately one thirtieth of the eventual full H-ATLAS sky coverage. In section \ref{Sample} we describe the sample that we have chosen to use for this analysis and the completeness corrections required. In section \ref{masses} we describe how we have derived luminosities and dust masses from the {\em Herschel} data, while in section \ref{DMF}, we present the dust mass function and evaluate its evolution. Section \ref{models} compares the DMF to models of dust evolution in order to explain the origin of the strong evolution. Throughout we use a cosmology with $\Omega_m = 0.27,\,\Omega_{\Lambda} =0.73$ and $H_o = 71\, \rm{km\,s^{-1}\,Mpc^{-1}}$.

\label{conclusions} We have estimated the dust mass function for the Science Demonstration Phase data from the {\em Herschel}-ATLAS survey, and investigated the evolution of the dust mass in galaxies over the past 5 billion years. We find that: \begin{itemize} \item{There is no evidence for evolution of dust temperature out to $z=0.5$ in this 250\mic selected sample.} \item{The dust mass function and dust mass density shows strong evolution out to $z=0.4-0.5$. In terms of pure mass evolution this corresponds to a factor 4--5 increase in the dust masses of the most massive galaxies over the past 5 billion years} \item{Similar strong evolution is found in the ratio of dust-to-stellar mass and V-band optical depth - {\em Herschel}-selected galaxies were more dusty and more obscured at $z=0.4$ compared to today.} \item{In order to account for the evolution of the dust content we need to radically alter chemical and dust evolution models. We cannot reproduce these trends with Milky Way metal or dust yields or star formation efficiencies.} \item{H-ATLAS 250$\mu$m selected sources are highly efficient at converting metals into dust, either through mantle growth or through a bias in the IMF towards higher mass stars. They must also be observed following an episode of star formation (either recent formation or recent major burst) where the gas has been consumed at a much faster rate than galaxies like the Milky Way today.} \item{As dust and gas (particularly molecular gas associated with SF) are tightly correlated in galaxies, this increase in dust content is suggestive of galaxies being more gas rich at $z=0.5$. According to the simple chemical model, we are possibly witnessing the period of growth toward peak dust mass when gas fractions are $\sim 0.5$ or higher. This strong decline in gas and dust content may be an explanation for the decrease in star-formation rate density in recent times as measured in many multi-wavelength surveys.} \end{itemize} This study uses only 3 percent of the area of the H-ATLAS data. Future improvements will come from the wider area coverage of the full survey, reducing uncertainties due to cosmic variance and small number statistics. Use of deeper optical/IR data from forthcoming surveys such as VISTA-VIKING, pan-STARRS, DES and VST-KIDS will also allow us to push to earlier times and higher redshifts to find the epoch of maximum dust content in the Universe.

1012.5186

1012

1012.0595_arXiv.txt

We perform simulations of general relativistic rotating stellar core collapse and compute the gravitational waves (GWs) emitted in the core bounce phase of three representative models via multiple techniques. The simplest technique, the quadrupole formula (QF), estimates the GW content in the spacetime from the mass quadrupole tensor only. It is strictly valid only in the weak-field and slow-motion approximation. For the first time, we apply GW extraction methods in core collapse that are fully curvature-based and valid for strongly radiating and highly relativistic sources. These techniques are not restricted to weak-field and slow-motion assumptions. We employ three extraction methods computing (i) the Newman-Penrose (NP) scalar $\Psi_4$, (ii) Regge-Wheeler-Zerilli-Moncrief (RWZM) master functions, and (iii) Cauchy-Characteristic Extraction (CCE) allowing for the extraction of GWs at future null infinity, where the spacetime is asymptotically flat and the GW content is unambiguously defined. The latter technique is the only one not suffering from residual gauge and finite-radius effects. All curvature-based methods suffer from strong non-linear drifts. We employ the fixed-frequency integration technique as a high-pass waveform filter. Using the CCE results as a benchmark, we find that finite-radius NP extraction yields results that agree nearly perfectly in phase, but differ in amplitude by $\sim 1-7\%$ at core bounce, depending on the model. RWZM waveforms, while in general agreeing in phase, contain spurious high-frequency noise of comparable amplitudes to those of the relatively weak GWs emitted in core collapse. We also find remarkably good agreement of the waveforms obtained from the QF with those obtained from CCE. The results from QF agree very well in phase and systematically underpredict peak amplitudes by $\sim5-11\%$, which is comparable to the NP results and is certainly within the uncertainties associated with core collapse physics.

\label{section:introduction} Massive stars ($M \gtrsim 8-10\,M_\odot$) end their nuclear burning lives with a core composed primarily of iron-group nuclei embedded in an onion-skin structure of progressively lighter elements. Energy generation has ceased in such a star's high-density core and relativistically-degenerate electrons provide pressure support against gravity. Silicon shell burning, neutrino cooling, and deleptonization eventually push the core over its effective Chandrasekhar mass. Radial instability sets in, leading to core collapse, accelerated by electron capture and photodisintegration of iron-group nuclei (see, e.g., \cite{bethe:90,baron:90}). The collapsing iron core separates into a subsonically collapsing homologous ($v\propto r$) inner core and supersonically infalling outer core. When the former reaches nuclear density, the nuclear equation of state (EOS) stiffens, dramatically increasing central pressure support and stabilizing the inner core, which, due to its large inertia, overshoots its new equilibrium and then rebounds into the still collapsing outer core, launching the hydrodynamic supernova shock. The acceleration experienced by the inner core in this \emph{core bounce} is tremendous, leading to the reversal of the collapse velocities of order $0.1 c$ of its $\sim 0.5\, M_\odot$ of material on a millisecond timescale. It was realized early on that the large accelerations encountered in stellar collapse in combination with a source of quadrupole (or higher) order asphericity lead to the emission of a burst of gravitational waves (GWs; see \cite{ott:09} for a historical overview). Rotation, centrifugally deforming the inner core to oblate shape, is an obvious source of such quadrupole asymmetry and rotating core collapse and bounce is the most extensively studied GW emission process in stellar collapse~(see, e.g., \cite{ott:07prl,dimmelmeier:08,scheidegger:10,scheidegger:10b,takiwaki:10,ott:11a} for recent studies and references therein). Alternatively, asymmetries in collapse may arise from perturbations, e.g., due to large convective plumes in the final phase of core nuclear burning, and may lead to GW emission at bounce and/or seed GW-emitting prompt postbounce convection~\cite{bh:96,fryer:04,ott:09}. A multitude of GW emission processes may be active in the postbounce, pre-explosion phase. These include convection/turbulence in the protoneutron star and in the postshock region, nonaxisymmetric rotational instabilities of the protoneutron star, protoneutron star pulsations, instabilities of the standing accretion shock, and asymmetric emission of neutrinos (\cite{ott:09,ott:06prl,marek:09,murphy:09,kotake:09,yakunin:10} and references therein). Of the entire ensemble of potential GW emission processes in stellar collapse, rotating core collapse and bounce is arguably the simplest and yields the cleanest signal, depending only on rotation, on the nuclear EOS, and on the mass of the inner core at bounce \cite{dimmelmeier:08}. Moreover, 3D studies have shown that collapsing iron cores with rotation rates in the range of what is physically plausible stay axisymmetric throughout the collapse phase and develop nonaxisymmetric dynamics only after bounce~\cite{shibata:05,ott:07prl,scheidegger:10}. Hence, the GW signal of rotating core collapse and bounce is linearly polarized and axisymmetric (2D) simulations are sufficient for its prediction. Unlike postbounce dynamics involving large scale and small scale fluid instabilities of stochastic nature, the GW signal of rotating collapse and bounce can, in principle, be predicted exactly for a given set of initial data. Hence, it has the potential of being used in GW searches using matched-filtering techniques (e.g., \cite{thorne:87}) or alternative approaches also taking into account detailed signal predictions~\cite{roever:09,summerscales:08}. Much progress has been made in recent years in the modeling of rotating core collapse and its GW signature. State-of-the-art simulations are general relativistic (GR) \cite{dimmelmeier:02,dimmelmeier:05,shibata:04,shibata:05, shibata:06,ott:07prl,ott:07cqg,kuroda:10,dimmelmeier:07, dimmelmeier:08} and some studies include magnetic fields \cite{shibata:06,obergaulinger:06b,kuroda:10} or finite-temperature EOS, deleptonization, and progenitors from stellar evolutionary calculations \cite{ott:07prl,ott:07cqg,dimmelmeier:07, dimmelmeier:08}. These improvements in the physics included in core collapse models provide for a more accurate and reliable dynamics underlying the emission of GWs. The calculation of the GW signal itself, however, is still being carried out predominantly in the slow-motion, weak-field quadrupole approximation (e.g., \cite{thorne:80}) that is of questionable quality, given the extreme densities and velocities involved in core collapse. The quadrupole formula (QF) ``extracts'' GWs based on matter dynamics alone, is not invariant under general relativistic gauge transformations, treats the emission region as a point source, and suffers from the fact that the definition of the generalized mass quadrupole moment is not unique in GR. In GR, the GW content of a spacetime can be extracted by means of the perturbative Regge-Wheeler-Zerilli-Moncrief (RWZM) formalism~\cite{Regge1957,Zerilli1970b,Zerilli:1971wd,Moncrief74} which is gauge invariant to first order or via the Newman-Penrose (NP) scalars approach~\cite{Newman1962,Penrose1963} which depends on the non-unique choice of the tetrad in which the Newman-Penrose scalars are evaluated. For reliable results, both RWZM and NP require extraction in the wave zone \cite{thorne:80} at coordinate radii many wavelengths from the source, but even there, coordinate ambiguities exist. The latter are removed only when GWs are extracted at future null infinity ($\scri^+$, see \cite{Newman1962,Penrose1963}), where space is asymptotically flat. Shibata \& Sekiguchi~\cite{shibatasekiguchi:03} have used simulations of an oscillating polytropic neutron star model to compare QF and finite-radius RWZM results. For the same basic system, Baiotti~et~al.~\cite{baiotti:09} compared QF, finite-radius RWZM, and finite-radius NP GW extraction with each other and with results from a 1D perturbation analysis. Both studies found that in the context of neutron star oscillations, the \emph{phase} of the waveforms obtained with the quadrupole approximation agrees exceptionally well with that of the RWZM and NP extraction methods. Shibata \& Sekiguchi, using their particular choice of the generalized quadrupole moment, found a systematic $\sim 20\%$ underprediction of the GW \emph{amplitudes} by the QF\@. Baiotti~et~al.~\cite{baiotti:09}, who studied multiple incarnations of the QF, found either underprediction or overprediction of the amplitude, both by up to $\sim60\%$, depending on the particular choice of QF\@. Nagar~et~al.~\cite{nagar:05b} studied the performance of RWZM and QF-based GW extraction from oscillating polytropic tori and found qualitatively similar results, and quantitative differences in amplitudes and integrated emitted energies $E_\mathrm{GW}$ between $\sim 2\%$ and $\sim 25\%$, again depending on the choice of quadrupole moment. RWZM and NP GW extraction and comparisons with the QF approximation for GWs emitted in core collapse spacetimes have proven difficult. On the one hand, the emitted GWs are weak: Typical strain amplitudes are $Dh \sim 10-1000\,\mathrm{cm}$, where D is the distance to the source, and typical emitted energies are of order $10^{-10}-10^{-8}\,M_\odot c^2$ \cite{ott:09}, many orders of magnitude lower than what is expected, for example, from double neutron star coalescence \cite{oechslin:07} or binary black hole mergers \cite{Reisswig:2009vc,Lousto:2009mf}. On the other hand, the GWs have typical frequencies of $100 - 1000\,\mathrm{Hz}$ and corresponding wavelengths of $300 - 3000\,\mathrm{km}$, hence require extraction at large coordinate radii where the grid resolution of core collapse simulations is typically too low to allow extraction of the relatively low-amplitude GWs emitted in core collapse (see, e.g., the discussion in \cite{ott:06phd}). Shibata \& Sekiguchi, in \cite{shibata:05}, were able to extract GWs with the RWZM formalism from an extreme core collapse model that developed a rotationally-induced large-scale nonaxisymmetric deformation after bounce, emitting GWs with $Dh \sim 20000\,\mathrm{cm}$. For this model, they found that the QF accurately predicts the GW phase, but underestimates the strain amplitude by $\sim 10\%$. Due to the aforementioned difficulties, these authors were unable to compare RWZM with QF for more moderate, axisymmetric models. Cerd\'a-Dur\'an~et~al.~\cite{cerda:05} performed core collapse simulations using a second-order post-Newtonian (2PN) extension of the conformal-flatness approximation to GR\@. Exploiting an approximate relationship of the non-conformal 2PN part of the metric to its GW part \cite{cerda:05}, they were able to extract GWs from their 2PN metric in standard axisymmetric rotating core collapse models. They found very close agreement (to a few percent in strain amplitude) between QF and 2PN GW signals for almost all considered collapse models. Siebel \textit{et al.}~\cite{Siebel03} performed nonrotating axisymmetric core collapse simulations by employing evolutions based on a fully general relativistic null cone formalism. They added nonspherical perturbations to the star, leading to the emission of GWs which they were able to extract with the Bondi news function at $\mathcal{J}^+$. Comparisons to the QF suggested a significant discrepancy in amplitude and frequency from the more reliable Bondi news result. The results of Shibata \& Sekiguchi~\cite{shibata:05} and of Cerd\'a-Dur\'an~et~al.~\cite{cerda:05} provide some handle on the performance of the QF approximation in core collapse spacetimes. The former study, while being performed in full GR, considered only a single extreme model. In addition, the authors were forced to extract GWs with RWZM at too small radii for completely reliable results. The latter study, while considering a broader ensemble of models, was restricted to 2PN without considering full GR, leaving room for doubts about the quality of their GW extraction technique. Finally, the results of Siebel \textit{et al.}~\cite{Siebel03} were limited to axisymmetry without rotation and are unreliable in the presence of strong shocks~\cite{Siebel03}. In this study, we readdress GW extraction from rotating core collapse spacetimes. We perform $3+1$ GR hydrodynamics simulations of rotating core collapse, for the first time in the core collapse context extracting GWs with RWZM, NP, and multiple QFs and comparing the results of these methods. In addition, and also for the first time in the present context, we utilize the Cauchy-Characteristic Extraction (CCE) approach~\cite{Winicour05, Bishop97b, Reisswig:2009us, Reisswig:2009rx, Babiuc:2010ze} that propagates the GW information to $\scri^+$ for completely gauge independent and unambiguous GW extraction. In choosing our models set, we are guided by Cerd\'a-Dur\'an~et~al.~\cite{cerda:05}, and draw precollapse configurations from the set of \cite{dimmelmeier:02}. These models are GR $n=3$-polytropic iron cores in rotational equilibrium and we evolve them with an analytic hybrid polytropic/$\Gamma$-law EOS used in many previous studies of rotating core collapse~\cite{zwerger:97,dimmelmeier:02, obergaulinger:06b,shibata:04,shibata:05,shibata:06,cerda:05}. For physically accurate GW signal predictions to be used in GW data analysis, a microphysically more complete treatment is warranted. Fortunately, recent results of studies employing such modeling technology (e.g., \cite{ott:07prl,ott:07cqg, dimmelmeier:07, dimmelmeier:08,abdikamalov:10}) show that, with a proper choice of EOS parameters, hybrid-EOS models are able to qualitatively and to some extent quantitatively reproduce the GW signals obtained with the much more complex and computationally intensive microphysical studies. Hence, for the purpose of this study, we resort to the simpler hybrid-EOS models. Our simulations employ the open-source {\tt Zelmani} GR core collapse simulation package~\cite{ott:09c} that is based on the {\tt Cactus Computational Toolkit} \cite{goodale:03, cactusweb} and the {\tt Einstein Toolkit} \cite{einsteintoolkitweb}. While using the full $3+1$ GR formalism, we limit our simulations to an octant of the 3D cube, using periodic boundary conditions on two of the inner faces of the octant and reflective boundary conditions on the third face. This limits 3D structure to even $\ell$ and $m$ that are multiples of $4$, which is not a limitation for the current study, since rotating core collapse and the very early postbounce evolution are likely to proceed nearly axisymmetrically~\cite{ott:07prl,scheidegger:08,scheidegger:10}. We note that, even though the GW signal in rotating core collapse is dominated by the $(\ell=2,m=0)$ '+' polarization mode, there is no reason to expect different behavior for other GW multipoles or polarizations and our results should translate to the non-axisymmetric case. The results of our simulations indicate that NP extraction yields results that agree well with those obtained from the most sophisticated CCE method. We observe differences in amplitude of $1-7\%$, depending on the model, while the agreement in phase is nearly perfect. We also find that the RWZM formalism yields unphysical high-frequency signal components that make this method less suitable for core collapse simulations where the signal is very weak. Finally, we note that the quadrupole approximation yields surprisingly close results to those obtained from CCE\@. While the phases nearly perfectly agree, the amplitude shows differences of $5-11\%$. This paper is structured as follows. In Sec.~\ref{sec:methods}, we discuss our methodology, initial data, and EOS details. Section~\ref{sec:waveextract} discusses the various GW extraction methods that we employ. In Sec.~\ref{sec:results}, we present our results and discuss them in detail. Finally, in Sec.~\ref{sec:summary}, we summarize and review our findings.

\label{sec:summary} We have performed a comparison study of four currently available GW extraction techniques in the context of axisymmetric rotating stellar core collapse. This study is the first to succeed in extracting GWs directly from axisymmetric core collapse spacetimes and the first to employ the fully coordinate independent CCE extraction method for non-vacuum spacetimes. We have performed core collapse simulations with simplified microphysics using a set of three representative initial configurations leading to GW signals of varying strength and signal morphology in quantitative agreement with what is expected from microphysically more complete models. In addition to having extracted waves with variants of the standard coordinate-dependent slow-motion, weak-field quadrupole formula, we have employed (i) the Regge-Wheeler-Zerilli-Moncrief (RWZM) formalism, (ii) extraction based on the Newman-Penrose (NP) scalar $\Psi_4$, and (iii) Cauchy-characteristic extraction (CCE\@). Of these three latter curvature-based methods, RWZM and NP extract GWs at a finite radius from the source, and hence, are generally prone to systematic errors arising from (i) near-zone effects, (ii) gauge ambiguities, and (iii) non-vanishing matter contributions. The CCE method, on the other hand, extracts waves gauge invariantly at future null infinity $\mathcal{J}^+$, that is, at an infinite distance from the source where gravitational radiation is unambiguously defined. Hence, it is subject only to small systematic errors due to the presence of matter fields at the CCE world-tube locations. An integral ingredient contributing to our success in extracting GWs from core collapse using curvature-based methods has been the removal of unphysical non-linear low-frequency drifts from the waveforms that otherwise would make a proper analysis largely impossible. This has been achieved by the application of fixed-frequency integration (FFI, \cite{Reisswig:2010di}) for time integration and filtering to yield the strain $h$. Comparing the waveforms obtained with the various extraction methods, we make a number of observations: (i) NP- and CCE-extracted waveforms converge with extraction and world-tube radius, respectively. The waveforms obtained with the RWZM formalism show spurious high-frequency components that no other method reproduces. A number of tests imply that the RWZM method may be less applicable to weak GW signals, at least at the currently accessible numerical resolutions and grid sizes. (ii) NP extraction, CCE, and even the quadrupole approximation, yield waveforms which agree well in phase, with differences in the time lags between successive peaks of $\lesssim0.05\,\rm{ms}$. Since the RWZM formalism is contaminated by unphysical high-frequency components, an accurate determination of the phasing compared to the other methods is largely impossible. (iii) The maximum amplitudes at core bounce are different by $\sim1-7\%$ in waveforms obtained with NP extraction and are systematically smaller by $\sim5-11\%$ in waveforms obtained with the QF compared to the waves obtained via CCE\@. Accordingly, CCE yields waveforms that result in slightly higher signal-to-noise-ratios (SNRs) ($\sim6-9\%$). (iv) Overall, the error of the waveforms computed with the quadrupole approximation are well within numerical errors and physical uncertainties. Unlike the waveforms obtained with the curvature-based methods, the quadrupole waveforms do not suffer from low-frequency drifts. In that respect, the quadrupole approximation is advantageous. We also observe that the quadrupole variant using ``physical`` velocity components \cite{dimmelmeier:02a} yields waves that are closer to those obtained via CCE\@. However, this finding may be true only for the core collapse case studied here and may not hold in general. (v) While it is unlikely that matched filtering approaches will be used in searches for GWs from core collapse in the near future, we have nevertheless computed GW template mismatches, a measure for the detectability of differences between waveforms. We find that when used in hypothetical matched-filtering GW searches, waveforms from NP extraction, CCE, and the QF would lead to the detection of the same model, while the waveforms computed with the RWZM formalism would generally not. There are two major drawbacks of our current work: (i) The curvature-based methods assume vacuum at the extraction spheres and world-tube locations. Hence, we must, in principle, extract at very large radii where the stress-energy tensor is zero. This, however, is currently not possible, since the collapsing star extends over the entire computational grid and larger grids are computationally prohibitive. (ii) All curvature-based methods yield waveforms with unphysical low-frequency drifts, requiring removal by spectral cut-off via FFI. This is particularly problematic in models with physical content below $\sim100\,\rm{Hz}$. A possible improvement of the low-frequency behavior could be achieved by the inclusion of matter terms in the CCE method, or alternatively, by enlarging the simulation domain such that the extraction takes place outside of the star and in pure vacuum. The latter could be efficiently achieved by employing multiblock techniques that cover the wavezone by a set of spherical grids \cite{zink:08b}. Finally, we point out that we have considered only the GW signal from rotating core collapse and bounce in this first study using curvature-based GW extraction from core collapse spacetimes. While our results may transfer to other GW emission processes in core collapse, this is by no means guaranteed. Further work will be needed to adress curvature-based GW extraction also from postbounce convection and the standing accretion shock instability, protoneutron star pulsations, rotational instabilities, and black hole formation.

1012.0595

1012

1012.3769_arXiv.txt

{ A successful implementation of thermal leptogenesis requires the re-heat temperature after inflation $T_R$ to exceed $\sim 2\times 10^9$~GeV. Such a high $T_R$ value typically leads to an overproduction of gravitinos in the early universe, which will cause conflicts, mainly with BBN constraints. Asaka and Yanagida (AY) have proposed that these two issues can be reconciled in the context of the Peccei-Quinn augmented MSSM (PQMSSM) if one adopts a mass hierarchy $m({\rm sparticle})>m({\rm gravitino})>m({\rm axino})$, with $m({\rm axino})\sim$~keV. In this case, sparticle decays bypass the gravitino, and decay more quickly to the axino LSP, thus avoiding the BBN constraints. In addition, thermally produced gravitinos decay inertly to axion+axino, also avoiding BBN constraints. We calculate the relic abundance of mixed axion/axino dark matter in the AY scenario, and investigate under what conditions a value of $T_R$ sufficient for thermal leptogenesis can be generated. A high value of PQ breaking scale $f_a$ is needed to suppress overproduction of axinos, while a small vacuum misalignment angle $\theta_i$ is needed to suppress overproduction of axions. The large value of $f_a$ results in late decaying neutralinos. We show that, to avoid BBN constraints, the AY scenario requires a rather low thermal abundance of neutralinos, while higher values of neutralino mass also help. We combine these constraint calculations along with entropy production from late decaying saxions, and find the saxion needs to be typically at least several times heavier than the gravitino. A successful implementation of the AY scenario suggests that LHC should discover a spectrum of SUSY particles consistent with weak scale supergravity; that the apparent neutralino abundance is low; that a possible axion detection signal (probably with $m_a$ in the sub-$\mu$eV range) exists, but no direct or indirect signals for WIMP dark matter should be observed. }

\label{sec:intro} A wide assortment of data from atmospheric, solar, reactor and accelerator experiments can all be explained in terms of massive neutrinos with large mixing angles which undergo flavor oscillations upon propagation through matter or the vacuum\cite{nu_review}. The emerging picture of the physics behind neutrino oscillation data is most elegantly explained by the presence of massive gauge singlet right-hand Majorana neutrino states $N_i$ ($i=1-3$ a generation index) which give rise to see-saw neutrino masses\cite{seesaw}: $m_{\nu_i}\simeq (f_{\nu_i}v)^2/M_{N_i}$ with $f_{\nu_i}$ the neutrino Yukawa coupling, $v$ the vev of the Higgs field, and $M_{N_i}\sim 10^9-10^{15}$~GeV. In addition to explaining neutrino oscillation data, the presence of massive $N_i$ states offers an elegant explanation of baryogenesis in terms of leptogenesis~\cite{lepto_review}, wherein the states $N_i$ exist in thermal equilibrium in the early universe, but decay asymmetrically to leptons versus anti-leptons. The lepton-anti-lepton asymmetry is then converted to a baryon asymmetry via $B$ and $L$ violating, but $B-L$ conserving, sphaleron effects\cite{krs}. To realize the thermal leptogenesis scenario, the lightest of the heavy neutrino masses $M_1$ must be $\agt 2\times 10^9$ GeV. In order to produce such states thermally, a re-heat temperature of the universe after inflation of $T_R\agt M_1 > 2\times 10^9$~GeV is required\cite{T_R}. Augmenting the Standard Model with a new, extremely high energy scale $M_{N_i}$ naturally leads to severe quadratic divergences in the Higgs sector which will need to be highly fine-tuned. The solution here is to also incorporate supersymmetry (SUSY), which reduces quadratic divergences to merely logarithmic, and ameliorates the fine-tuning problem\cite{wss}. While the addition of weak scale softly broken SUSY into the theory is actually supported by the measured values of the gauge couplings from LEP experiments, it also leads to new conundrums such as the gravitino problem: the production of gravitinos in the early universe can lead to {\it (i)}~overproduction of LSP dark matter ({\it e.g.} the lightest neutralino) beyond relic density limits obtained from WMAP and other experiments, or {\it (ii)}~disruption of the successful explanation of Big Bang nucleosynthesis by introducing late decaying quasi-stable particles whose decay products can break up the newly minted light elements. The common solution to the gravitino problem\cite{gravprob} is to require a sufficiently low re-heat temperature such that thermal gravitino production is suppressed enough to avoid overproduction of dark matter or disruption of BBN\cite{kl}. For gravitino masses in the few TeV or below range, a value of $T_R\alt 10^5$~GeV is required. Naively, this is in obvious conflict with the $T_R$ requirements of thermal leptogenesis. A variety of solutions have been proposed to reconcile leptogenesis with the SUSY gravitino problem. One is to abandon the ``thermal'' aspect of leptogenesis, and invoke non-thermal leptogenesis wherein the heavy neutrino states are produced via some scalar field decay, for instance the inflaton\cite{ntlepto}. Another suggestion is to invoke the gravitino as LSP, so it does not decay. However, the gravitino LSP scenarios fall back into the BBN problem since then the NLSP SUSY particle suffers a late decay into gravitino plus SM states which again injects high energy particles into the post-BBN plasma. One solution is to speed up NLSP decay via a small component of $R$-parity violation\cite{covi,covi2}. In a recent work\cite{Baer:2010kw}, we proposed an alternative scenario, invoking mixed axion/axino dark matter, as would occur in the Peccei-Quinn\cite{pq,ww,ksvz,dfsz} augmented MSSM (the PQMSSM)\cite{pqmssm,axino}. In this case, we invoked models with very heavy gravitinos, $m_{\tG}\agt 10$~TeV, so that gravitinos decay before the onset of BBN. Then, overproduction of dark matter can be avoided by requiring an axino LSP with mass $m_{\ta}\sim 0.1-1$~MeV. Neutralinos produced either thermally or via gravitino decay will themselves decay typically to states such as $\ta\gamma$, so that the dark matter abundance is reduced by a factor $m_{\ta}/m_{\tz_1}$\cite{ckr}. The bulk of dark matter then resides in thermally produced axinos and/or in axions produced from vacuum mis-alignment. By driving up the value of PQ breaking scale $f_a/N$, thermal production of axinos is suppressed, and higher values of $T_R$ are required to maintain a total axino plus axion relic abundance of $\Omega_{a\ta}h^2\sim 0.11$. To avoid overproduction of axions at high $f_a/N$, we adopted a small vacuum mis-alignment angle $\theta_i\sim 0.05$. However, the large values of $f_a/N\sim 10^{12}-10^{13}$~GeV suppress the $\tz_1$ decay rate, thus interfering with BBN from a different avenue. Models with a high-mass, bino-like $\tz_1$ and low ``apparent'' $\Omega_{\tz_1}^{app}h^2$ can avoid the BBN bounds, and allow $T_R$ values in excess of $10^{10}$~GeV to be attained. As we showed, such conditions with $m_{\tG}\sim 10-30$~TeV can be reached in Effective SUSY\cite{ckn,esusy} or mirage-unification SUSY breaking\cite{MU} models. A related scenario to reconcile thermal leptogenesis with the gravitino problem --- using mixed axion/axino dark matter --- was proposed much earlier by Asaka and Yanagida (AY)\cite{ay}. Their proposal was to work within the context of the PQMSSM, but with a sparticle mass hierarchy $m({\rm sparticle})>m_{\tG}>m_{\ta}$. In this case, the couplings of MSSM sparticles to axinos are larger than the couplings to gravitinos, so that the long-lived decays to gravitino are bypassed, and the sparticles are assumed to decay to an axino LSP shortly before the onset of BBN. Furthermore, thermally produced gravitinos decay inertly via $\tG\to a\ta$ and so do not disrupt BBN. Reheat temperatures as high as $T_R\sim 10^{15}$ were claimed to be possible. In this paper, we re-visit the AY scenario, incorporating several improvements into our analysis. In particular, we implement \begin{enumerate} \item the latest astrophysically measured value of\cite{wmap7} \be \Omega_{\rm DM}h^2= 0.1123\pm 0.0035\ \ \ {\rm at\ 68\%\ CL}; \ee \item the latest calculations for thermal production of gravitinos and axinos; \item vacuum-misalignment production of axions as an element of the dark matter abundance; \item the latest BBN constraints on late decaying particles; and finally \item a careful treatment of entropy production from late decaying saxions. Since entropy production from saxion decay will also dilute the matter-antimatter asymmetry by a factor $r$ (to be defined later), in this case a re-heat temperature $T_R\agt 2r\times 10^9$ GeV will be needed. \end{enumerate} The re-analysis of the AY scenario taking into account points 1.--4.\ is presented in Sec.~\ref{sec:AY}, while entropy injection from saxion decay is discussed in detail in Sec.~\ref{sec:saxion}. In Sec.~\ref{sec:conclude}, we present our final conclusions and consequences of the AY scenario for LHC physics and dark matter searches.

\label{sec:conclude} In this paper, we reported on investigations of the viability of the Asaka-Yanagida suggestion that a mass hierarchy with $m({\rm sparticle})>m_{\tG}>m_{\ta}$ can be used to reconcile thermal leptogenesis, which requires $T_R\agt 2\times 10^9$~GeV, with the gravitino problem, which seemingly requires much lower $T_R$ to avoid BBN constraints and overproduction of neutralino dark matter. In the AY scenario, the $\tG$ decays inertly to $a\ta$. BBN constraints on ${\rm sparticle}\to \tG +{\rm particle}$ can be avoided because the much faster decays ${\rm sparticle}\to \ta + {\rm particle}$ are now allowed. We re-examined the AY scenario in Sec. \ref{sec:AY} by including 1.~updated measurements on the total dark matter abundance $\Omega_{DM}h^2\simeq 0.1123$, 2.~updated calculations of thermal axino and gravitino production, 3.~the contribution of relic axions and 4.~BBN constraints on late decaying $\tz_1$s. Furthermore, in Sec.~\ref{sec:saxion}, we included dilution of dark matter by saxion production and decay. The latter effect can be neglected if $m_s$ is in the multi-TeV range and the initial saxion field strength $s_i$ is somewhat smaller than the PQ breaking scale $f_a/N$. We found in Sec.~\ref{sec:AY}, neglecting the saxion entropy effect, that the AY scenario does work under the conditions that {\it (i)}~$f_a/N$ is rather large $\agt 10^{12}$~GeV, implying a somewhat lighter axion than is presently searched for by ADMX\cite{admx}, {\it (ii)}~the apparent neutralino relic density $\Omega_{\tz_1} h^2$ is not too big: $\Omega_{\tz_1} h^2\alt 1$, {\it (iii)}~the value of $m_{\tz_1}$ is at least in the several hundred GeV range in order to hasten the $\tz_1$ decay rate, and {\it (iv)}~the axion mis-alignment angle $\theta_i$ is on the small side $\alt 0.5$ to suppress overproduction of axions when $f_a/N$ is large. By including saxion production and decay in Sec. \ref{sec:saxion}, we can dilute the axino and also axion DM abundance, which in turn allows for somewhat higher values of $T_R$ up to $\sim 10^{13}$~GeV to be generated. However, since saxion decay also dilutes the baryon density, in this case we must require instead $T_R/r> 2\times 10^9$ GeV. The saxion mass $m_s$ needs to be rather large to avoid BBN constraints on late decaying saxions if $T_R$ is to be high. In this case, the DM is likely to be {\it mainly axions}, although a few cases with mainly axino DM were generated. The axion mis-alignment angle need not be small here since the axion abundance can be suppressed by entropy injection from saxions. We have also found that a large portion of the MSSM parameter space ($\Omega_{\tz_1}$ and $m_{\tz_1}$) can be consistent with high $T_R$ and still avoid the BBN bounds on late decaying neutralinos, due to the dilution of the neutralino relic density after the entropy injection from saxion decays. The observable consequences of our final results are as follows. If the AY scenario with $m({\rm sparticle})>m_{\tG}>m_{\ta}$ is to reconcile thermal leptogenesis with the gravitino problem, then we expect several broad results to ensue: \begin{enumerate} \item discovery of SUSY at the LHC, with a reconstructed $\Omega_{\tz_1} h^2$ not too large, lest $\tz_1$s are produced at too large a rate in the early universe, and their late decays disrupt BBN; \item a SUSY mass spectrum consistent with SUGRA models with a rather light (but still weak scale) gravitino, since the gravitino mass must be lighter than all observable sparticles; \item a mainly bino-like $\tz_1$, to quicken decays into $\ta\gamma/Z$, with mass $m_{\tz_1}$ in the hundreds of GeV range, which also helps diminish the lifetime; \item no direct or indirect detection of neutralino (WIMP) dark matter; \item finally, we expect discovery of an axion to be likely, but in the mass range $\sim 0.1-2$ $\mu$eV, somewhat below the values presently being explored. \end{enumerate}

1012.3769

1012

1012.4655_arXiv.txt

The heating of solar chromospheric inter-network regions by means of the absorption of electromagnetic (EM) waves that originate from the photospheric blackbody radiation is studied in the framework of a plasma slab model. The absorption is provided by the electron-neutral collisions in which electrons oscillate in the EM wave field and electron-neutral collisions damp the EM wave. Given the uncertain nature of the collision cross-section due to the plasma micro-turbulence, it is shown that for plausible physical parameters, the heating flux produced by the absorption of EM waves in the chromosphere is between $20 - 45$ \% of the chromospheric radiative loss flux requirement. It is also established that there is an optimal value for the collision cross-section, $5 \times 10^{-18}$ m$^{2}$, that produces the maximal heating flux of 1990 W m$^{-2}$.

The problem of heating of the solar atmosphere, the sharp temperature raise from photospheric 6000K to few $10^6$K in the corona has been a long standing problem of solar physics. There is no lack of possible heating mechanisms in the {\it corona} with the two main candidates being so-called direct current (DC) models that are based on the magnetic reconnection and alternating current (AC) models that are based on magnetohydrodynamic (MHD) wave dissipation \cite{2005psci.book.....A}, alongside with few dozen other, less widely accepted \cite{1994ApJ...427..446S,2006A&A...455.1073T} or less successful ones \cite{2005A&A...441.1177T}. In the {\it chromosphere}, however, until recently, a general agreement was that it is heated by the absorption of the acoustic waves. A distinction needs to be drawn between wave heating of different parts of the chromosphere. The chromosphere of the quiet Sun can be broadly split into two parts: (i) the magnetic network, which marks the boundaries in between the super-granulation cells; and (ii) the inter-network regions, which constitute the bulk surface area of the chromosphere (i.e. the super-granulation cell interiors). In the magnetic network the magnetic field is nearly radial (vertical) and quite strong (of the order of few kG). Since the strong magnetic field there provides a substantial amount of free energy, in principle it seems reasonable to believe that the magnetic network can be heated by the dissipation of MHD waves \cite{2008ApJ...680.1542H}. However, the role of magnetic reconnection in the heating of the magnetic network cannot be discounted. For example, there seems to be an evidence of forced magnetic reconnection taking place in the photosphere \cite{2010ApJ...712L.111J}, as well as in chromosphere \cite{2010ApJ...713L...6C}. The estimates of the heating flux produced by plausible reconnection models seem to fall short by up to two orders of magnitude from the quiet chromosphere and coronal heating requirements \cite{1999ApJ...524..483L}. At the same time Ref.\cite{2008ApJ...680.1542H} and other similar works, that base their conclusions on the numerical simulation results, do not make precise predictions for the heating rates produced by the dissipation of MHD waves in the chromospheric magnetic network. The situation with the heating of inter-network regions where magnetic fields are weak is even worse than with the magnetic network. On one hand, this is because of the lack of the magnetic energy. On the other hand, the results of Ref.\cite{2005Natur.435..919F} indicate that the acoustic energy flux of the high-frequency (10-50mHz) acoustic waves (that were previously believed to constitute the dominant heating mechanism of the chromosphere) falls short, by a factor of at least ten, to balance the radiative losses in the solar chromosphere. This led them to a conclusion that the acoustic waves cannot constitute the dominant heating mechanism of the solar chromosphere. This conclusion has been challenged by Ref.\cite{2008JApA...29..163K}, who suggests that the observations reported by Ref.\cite{2005Natur.435..919F} only detect 10\% of the acoustic wave flux perhaps because of the limited spatial resolution. Ref.\cite{2010ApJ...723L.134B} report somewhat higher than usual flux carried by the acoustic waves at photospheric heights (250 km) as well as compile a useful list of up to date acoustic heating flux measurements. In this context, in this paper we explore an alternative to the acoustic wave heating idea of the quiet chromosphere. In particular, we investigate the following proposition: (i) It is known that the solar irradiance spectrum, that comes out of photosphere, is well approximated by an effective blackbody at a temperature of $T=5762$ K, in the frequency range of $f =30-1667$ THz (corresponding to the wavelengths range of $10 - 0.18$ $\mu$m) (see e.g. Figure 2.3 from Ref.\cite{2005psci.book.....A}). Therefore, we assume that the radiative heating flux with the Planckian brightness distribution as a function of frequency penetrates the lower part of the solar atmosphere (photosphere, $h=0 - 500$ km and chromosphere, $h=500-2200$ km). (ii) Instead of solving radiative transfer equations, we take the photospheric blackbody flux of $T=5762$ K, and quantify how much electromagnetic (EM) radiative flux is absorbed using a plausible model for EM wave absorption, which is based on Ref.\cite{2003ITPS...31..405T} plasma slab model combined with the VAL-C model of chromosphere \cite{1981ApJS...45..635V}. Ref.\cite{2003ITPS...31..405T} plasma slab model is based on splitting a smoothly varying, non-uniform density, weakly ionised plasma with the uniform magnetic field along the density gradient, into a set of thin sub-slabs with a uniform density in each sub-slab - thus providing a discretized version of the smooth density profile. The absorption of the EM radiation is based on two physical effects: electron-neutral collisions and electron cyclotron resonance. For the considered in our model radial magnetic field value of 0.2 kG, electron cyclotron frequency is $f_{ce}=eB/2 \pi m_e= 0.00056$ THz. Also, for the considered model parameters, the ratio of electron-neutral collision frequency and electron cyclotron frequency, $\nu_{en}/f_{ce} << 1$, which ensures that the collisions would not affect any electron cyclotron resonance damping. Therefore EM wave absorption via electron cyclotron resonance is negligibly small in the considered range of frequencies $2 - 2000$ THz. We refer the interested reader to Ref.\cite{2003ITPS...31..405T} for the details of the plasma slab model. However, we re-iterate the key points of the model in Section 2. As a result we find that for plausible physical parameters, the heating flux produced by the absorption of EM waves in the chromosphere is between $20 - 45$ \% of the VAL-C radiative loss flux requirement. We also establish that there is an optimal value for the collision cross-section, $5 \times 10^{-18}$ m$^{2}$, that produces the maximal heating flux of 1990 W m$^{-2}$. The paper is organised as follows: In Section 2 we describe the model. In Section 3 we present the numerical results and we close with the Conclusions in Section 4.

In this paper we put to test a simple proposition that some of the problems of the chromospheric inter-network regions (regions with weak magnetic field which constitute the bulk of solar chromosphere surface), discussed in the Introduction section can be alleviated by inclusion of the absorption of photospheric EM radiation in the plasma sub-slab based model. On one hand, we know that that the solar irradiance spectrum, that comes out of photosphere, is well approximated by an effective blackbody at a temperature of $T=5762$ K, in the frequency range of $f =30-1667$ THz. Therefore, we can assume that the radiative heating flux with the Planckian brightness distribution as a function of frequency illuminates lower part of the solar atmosphere (photosphere, $h=0 - 500$ km and chromosphere, $h=500-2200$ km). On the other hand, instead of solving radiative transfer equations, we can take photospheric blackbody flux of $T=5762$ K, and quantify how much electromagnetic (EM) radiative flux is absorbed using a plausible model for the EM wave absorption which is based on Ref.\cite{2003ITPS...31..405T} plasma slab model combined with VAL-C model of chromosphere \cite{1981ApJS...45..635V}. Our model is based on splitting a weakly ionised plasma slab with the uniform magnetic field along a smooth density gradient, into a set of narrow sub-slabs with a uniform density in each slab - hence providing a discretized version of the smooth density gradient. In the relevant frequency range ($2 - 2000$ THz), the absorption of the EM radiation is due to the electron-neutral collisions, while the electron cyclotron resonance can be ignored. The absorption of the EM radiation due to the electron-neutral collisions happens because the electrons oscillate in the EM wave field and electron-neutral collisions (i.e. their mutual friction) then results in the damping of the EM wave. We also include a contribution to the cross-section from the anomalous plasma micro-turbulence, which we incorporate in an additive way. Our model has two potential weaknesses: (i) to what extent the radiative heating flux of the photosphere deviates from the Planckian brightness distribution as a function of frequency? and (ii) whether the {\it absorption} of EM radiation can be described by the plasma sub-slab model which assumes local thermodynamic equilibrium (LTE)? Concerning the first question we state that it is well known that the solar irradiance spectrum, that comes out of photosphere, is well approximated by an effective blackbody at a temperature of $T=5762$ K, in the frequency range of $f =30-1667$ THz. Also, there are plausible models of solar photosphere that use LTE assumption (see e.g. Ref.\cite{2008A&A...486..951G}) that implies the applicability of the Planckian brightness function. As to the second question we remark that indeed in order to properly work out the chromospheric (radiative) {\it losses}, one needs to consider non-LTE effects. Strong radiation field in the chromosphere can drastically alter the occupation numbers of the energy levels in atoms, thus producing non-LTE net radiative loss \cite{1978stat.book.....M}. We assert, however, that the chromospheric {\it heating} process can be regarded as an LTE process, given the {\it steady} inflow of EM radiation from the photosphere. After all, the chromospheric heating models that are based on the absorption of acoustic shock wave energy use equations of hydrodynamics which imply LTE, and it is only the radiative {\it losses} that are treated by the non-LTE radiative transport as in e.g. Ref.\cite{2005Natur.435..919F}. As a result we find that: (i) for plausible physical parameters, the heating flux produced by the absorption of EM waves in the chromosphere is between $20 - 45$ \% of the VAL-C model radiative loss flux requirement. The variation range is because of the uncertainties in the collision cross-section due to the plasma micro-turbulence. (ii) We also established that for absorption in the region 500 km - 2200 km above photosphere there is an optimal value for $\sigma = 5 \times 10^{-18}$ m$^{2}$ that produces the maximal heating flux of 1990 W m$^{-2}$. From the observational point of view, if the absorption of EM waves in the frequency range $2 - 2000$ THz has a significant role to play in the heating of quiet chromosphere, as suggested by our findings, then the plasma slab model also predicts that: (i) There is a good case for the electron-neutral anomalous collision cross-section to be a factor of 10 larger than the value predicted by the classical plasma transport. Ref.\cite{2003PhPl...10..319T} presented plasma resistivity (which is proportional to both the collision frequency and cross-section) measurements in the reconnection current sheet of the Magnetic Reconnection Experiment. They established that in some regimes, the measured resistivity values can be more than an order of magnitude larger than the classical Spitzer value. Therefore it would seem likely that the collision cross-section in the chromosphere would also assume some anomalous value. (ii) There should be a good correlation of the total solar irradiance with the Mg-index (which represents the chromospheric excess radiation relative to the photosphere) on a long-term (1 month or more) timescale. This is because our model takes photospheric blackbody EM wave flux as the source of energy, that irradiates chromosphere from below (a torch shining from the below analogy is relevant here). In fact, this is exactly what is observed: Ref.\cite{2007ASPC..368..481S} presents the data that shows that the total solar irradiance and the Mg-index have a correlation coefficient of 0.8 using monthly data averages (see their Figure 1 and pertinent discussion). The correlation is somewhat worse of a shorter timescales, e.g. daily data averages -- this has a good explanation in that contribution from the solar features such as sunspots and faculae (that affect photospheric total solar irradiance) and plages (that affect chromospheric brightness and are in fact mapped closely to the faculae below) average out on the long timescales and generally track to solar activity cycle (that has a proxy of number of sunspots). (iii) Unlike in the photosphere, the chromospheric brightness should not decrease with the increase of the magnetic field. This is because in our model EM wave absorption depends on the magnetic field rather weakly. This is also what is observed: Ref.\cite{1989ApJ...337..964S} find that CaII K line core contrast (the relative difference between the intensity at a given magnetogram and the quiet Sun intensity) that is a measure of chromospheric brightness is weakly increasing with the magnetic field as $\propto B^{0.6}$. In the photosphere the contrast of a continuum in the green part of the solar spectrum initially increases with $B$ up to 0.02 Tesla but then sharply decreases with $B$ above 0.05 Tesla. Ref.\cite{2007ASPC..368..481S} explain chromospheric rise of brightness with the increase of magnetic field for small $B$'s ($0.01$ T) is due to the density increase of the magnetic flux tubes, and for large $B$'s ($> 0.05$T) subsequent slower rise is due to the quenching of the wave activity. As Ref.\cite{1989ApJ...336..475T} have shown, the strong magnetic fields inhibit average horizontal flow speeds in the granules. Thus in conclusion the plasma slab model predictions seem also to conform with the available observational data.

1012.4655

1012

1012.1246_arXiv.txt

The orbital evolution of a dust particle under the action of a fast interstellar gas flow is investigated. The secular time derivatives of Keplerian orbital elements and the radial, transversal, and normal components of the gas flow velocity vector at the pericentre of the particle's orbit are derived. The secular time derivatives of the semi-major axis, eccentricity, and of the radial, transversal, and normal components of the gas flow velocity vector at the pericentre of the particle's orbit constitute a system of equations that determines the evolution of the particle's orbit in space with respect to the gas flow velocity vector. This system of differential equations can be easily solved analytically. From the solution of the system we found the evolution of the Keplerian orbital elements in the special case when the orbital elements are determined with respect to a plane perpendicular to the gas flow velocity vector. Transformation of the Keplerian orbital elements determined for this special case into orbital elements determined with respect to an arbitrary oriented plane is presented. The orbital elements of the dust particle change periodically with a constant oscillation period or remain constant. Planar, perpendicular and stationary solutions are discussed. The applicability of this solution in the Solar system is also investigated. We consider icy particles with radii from 1 to 10 $\mu$m. The presented solution is valid for these particles in orbits with semi-major axes from 200 to 3000 AU and eccentricities smaller than 0.8, approximately. The oscillation periods for these orbits range from 10$^{5}$ to 2 $\times$ 10$^{6}$ years, approximately.

\label{sec:introduction} Since \cite{poynting} and \cite{robertson}, the dynamics of dust grains has been investigated in many papers and its analysis is an inseparable part of astrophysics. The influence of the electromagnetic radiation of a central star is usually taken into account in the form of the Poynting-Robertson (P-R) effect (\citealt{poynting,robertson,wywh,burns,klacka}). Beside the P-R effect, the stellar wind (corpuscular radiation of the star) can also affect the motion of dust particles. Covariant derivation of the acceleration caused by the stellar wind and its effects on the dynamics of dust grains were presented in \cite{SW}. The Solar magnetic field can affect the motion of charged dust particles in the Solar system (\citealt{parker,kimura}). Planets can capture the dust particles into the mean-motion resonances (\citealt{dermott,reach}). Because of the relative motion of the Sun with respect to the local interstellar medium, atoms of the interstellar medium approach the Sun. The direction of the approach is given by the actual velocity of the Sun with respect to the local interstellar medium. These approaching atoms form an interstellar gas flow. This flow of interstellar atoms through the Solar system has been already investigated in the past (e.g. \citealt{fahr,mobius,alouani}). This interstellar gas flow can affect the dynamics of dust grains in outer parts of the Solar system. The assumption that dust particles in orbits around other stars can be also affected by interstellar gas flow is supported by the recent detection of debris disks around stars with asymmetric morphology caused by a fast motion of the disk through a cloud of interstellar matter (\citealt{hines,debes}). The influence of interstellar gas flow on the dynamics of a spherical dust particle was investigated by \cite{scherer}, who calculated the secular time derivatives of the particle's angular momentum and the Laplace-Runge-Lenz vector caused by the interstellar gas flow. He has come to the conclusion that the particle's semi-major axis can increase exponentially (\citealt[p. 334]{scherer}). This result contradicts results of \cite{flow}. Here the secular time derivatives of the Keplerian orbital elements of the dust particle under the action of a fast interstellar gas flow were for the first time calculated for arbitrary orientations of the orbit. \cite{flow} states that the secular semi-major axis of the dust particle must decrease under the action of the interstellar gas flow. The rate of decrease is proportional to the semi-major axis. Decrease of the semi-major axis only happens in secular time-scales. Therefore, the exact calculation of the secular time derivative of the semi-major axis was necessary. This result was confirmed also by \cite{bera} who investigated the motion of a dust particle in the outer region of the Solar system behind the solar wind termination shock. They calculated the orbital evolution of the dust particle under the action of a constant mono-directional force, i.e., they solved a case of the classical Stark problem (for more information about the Stark problem see \citealt{lantoine} and the references therein). In the Stark problem it is assumed that the orbital speed of the dust grain with respect to a central object can be neglected in comparison with the speed of the interstellar gas flow. The relative velocity terms are in \cite{flow} taken into account to the first order of accuracy in the calculation of the secular time derivatives of Keplerian orbital elements. \cite{bera} reproduced the result of \cite{flow} on the secular evolution of the semi-major axis of the particle's orbit. However, the case of a mono-directional force (i.e., the Stark problem) can be important for those stars where the relative speed of the neutral gas with respect to the central star is high. Such a situation can occur, for example, in merging galaxies when a star from the first galaxy moves through a molecular cloud of the second galaxy. \cite{bera} used the orbit-averaged Hamiltonian method. In this paper we use a method based on orbit-averaged time derivatives of Keplerian orbital elements. Using this different method we confirm results of \cite{bera}. We not only use a new method, we also find new results. We find an explicit form for the time dependence of all Keplerian orbital elements determined with respect to a reference plane perpendicular to the gas flow velocity vector. Mainly the time dependence of the longitude of the ascending node will represent a generalization of \cite{bera} results. We find a transformation for the orbital elements determined with respect to the reference plane perpendicular to the gas flow velocity vector into an arbitrarily oriented reference plane. We determine the maximal and minimal values of eccentricity for numerous special cases. We study in detail the behaviour of the solution for the planar case and for the case when the gas flow velocity vector is perpendicular to the line of apsides. Properties of the solution in the Solar system are discussed.

\label{sec:conclusion} We have investigated the long-term orbital evolution of a spherical dust particle perturbed by a small constant mono-directional force caused by a fast interstellar gas flow. The secular time derivatives of the particle's Keplerian orbital elements were derived. We transformed the system of differential equations for the Keplerian orbital elements into a system of differential equations for the semi-major axis, eccentricity, and the radial, transversal, and normal components of the interstellar gas flow velocity vector determined at the pericentre of the particle's orbit. In these new variables, the system of differential equations can be easily solved. The solution of the system was used in order to obtain the evolution of the Keplerian orbital elements in a reference frame in which the orbital elements are determined with respect to the plane perpendicular to the interstellar gas velocity vector. We found an explicit form for the time dependence of all Keplerian orbital elements in this reference frame. We generalized the expression for the time dependence of the longitude of the ascending node found by \cite{bera}. We transformed newly found orbital elements into an arbitrary reference frame. This transformation gave us explicit time dependences of all Keplerian orbital elements in the arbitrary reference frame. The orbital elements of the dust particle in an arbitrary reference frame change periodically with a constant oscillation period or else remain constant. We determined the properties of the solution for the planar case and for the case in which the gas flow velocity vector is perpendicular to the line of apsides. We found the maximal and minimal values of the eccentricity in these cases. In the planar case, the particle's orbit approaches the direction with maximal value of $I$. In the perpendicular case, the orbital plane is tilted back and forth around the line of apsides. We also confirmed the stationary solution found by \cite{bera}. For the stationary solution, the orbital plane rotates around the line aligned with the gas flow velocity vector and going through the centre of gravity. This solution can be applied also for the dust particles in the Solar system. If we consider icy particles with radii from 1 to 10 $\mu$m, then the solution is valid for orbits with semi-major axis from 200 to 3000 AU, approximately. More exact values of these limits depend on the radius of the particle. A maximal orbit eccentricity for which the solution is valid is smaller than approximately 0.8 and depends on the semi-major axis of the orbit and the particle's radius (see Fig. \ref{F3}). The period of change of the orbital elements for these orbits ranges from 10$^{5}$ to 2 $\times$ 10$^{6}$ years, approximately. \appendix

1012.1246

1012

1012.3899_arXiv.txt

Interferometric millimeter observations of the cosmic microwave background and clusters of galaxies with arcmin resolutions require antenna arrays with short spacings. Having all antennas co-mounted on a single steerable platform sets limits to the overall weight. A 25~kg lightweight novel carbon-fiber design for a 1.2~m diameter Cassegrain antenna is presented. The finite element analysis predicts excellent structural behavior under gravity, wind and thermal load. The primary and secondary mirror surfaces are aluminum coated with a thin TiO$_2$ top layer for protection. A low beam sidelobe level is achieved with a Gaussian feed illumination pattern with edge taper, designed based on feedhorn antenna simulations and verified in a far field beam pattern measurement. A shielding baffle reduces inter-antenna coupling to below $\sim$ -135~dB. The overall antenna efficiency, including a series of efficiency factors, is estimated to be around 60\%, with major losses coming from the feed spillover and secondary blocking. With this new antenna, a detection rate of about 50 clusters per year is anticipated in a 13-element array operation.

The Array for Microwave Background Anisotropy (AMiBA) is a forefront radio interferometer for research in cosmology. This project is led, designed, constructed, and operated by the Academia Sinica, Institute of Astronomy and Astrophysics (ASIAA), Taiwan, with major collaborations with National Taiwan University, Physics Department (NTUP), Electrical Engineering Department (NTUEE), and the Australian Telescope National Facility (ATNF). Contributions also came from the Carnegie Mellon University (CMU), and the National Radio Astronomy Observatory (NRAO). As a dual-channel 86-102 GHz interferometer array of up to 19 elements, AMiBA is designed to have full polarization capabilities, sampling structures on the sky greater than 2 arcmin in size. The AMiBA target science is the distribution of high red-shift clusters of galaxies via the Sunyaev-Zel'dovich Effect (SZE), e.g. \citet{sz72,bi99,carl02} and references therein, as a means to probe the primordial and early structure of the Universe. AMiBA will also measure the Cosmic Microwave Background (CMB), e.g. \citet{spergel07,aghanim08,larson10}, temperature anisotropies on scales, which are sensitive to structure formation scenarios of the Universe. AMiBA is sited on Mauna Loa in Hawaii, at an elevation of 3,400m to take advantage of higher atmospheric transparency and minimum radio frequency interference. After an initial phase with seven 0.6~m diameter antennas \citep{koch06} in a compact configuration, the AMiBA is currently operating with 13 1.2~m diameter Cassegrain antennas (Figure \ref{front}). This new antenna and its capabilities are described here. Section \ref{design} lists the antenna requirements. In Section \ref{mechanical} the mechanical and optical designs are detailed out, including simulation results of the structure and the antenna-feedhorn system. Section \ref{measurement} is devoted to the antenna verification measurements. The factors composing the antenna efficiency are estimated in Section \ref{efficiency}. Section \ref{novelty} discusses the improved design features of the 1.2~m antenna and the upgraded array operation. Our conclusion is given in Section \ref{conclusion}. Previous AMiBA progress reports were given in \citet{ho04,raffin04,li06,raffin06}. A project overview is given in \citet{ho08}. More details about the correlator and receiver can be found in \citet{li10} and \citet{chen08}. The hexapod telescope mount is introduced in \citet{koch08}. Observing strategy, calibration scheme and data analysis with quality checks are described in \citet{lin08,wu08,nishioka08}. First AMiBA science results from the 7-element array are presented in \citet{huang10,liao10,koch08b,liu08,umetsu08,wu08}. A possible science case utilizing the 1.2~m antenna array configuration is outlined in \citet{molnar10}.

\label{conclusion} A 1.2~m $f/0.35$ Cassegrain antenna for a single platform close-packed interferometer for astronomical radio observations is presented. Due to weight constraints, carbon fiber reinforced plastic (CFRP) is chosen as a lightweight material for the main antenna parts. With a detailed finite element analysis (FEA) it has been possible to keep the weight within 25~kg. The primary and secondary mirror sandwich composite structures show excellent behavior under thermal, wind and gravity load, leading to FEA predicted surface rms deformation errors of less than 10~$\mu$m and maximum tilts in the optical axis of about 1 arcmin. The primary paraboloid and secondary hyperboloid mirror manufactured surface rms errors are typically around 30~$\mu$m and 10~$\mu$m, respectively. The mechanical alignment after shimming and the resulting focal length are within $\sim 0.1$~mm of the specifications. The efficiency loss due to mechanical assembly and manufacturing is then within $\sim$ 1\%. For a good reflectivity the mirrors are coated with a $\sim2\mu$m aluminum layer. Additionally, on top of that, a thin TiO$_2$ layer ($\sim 0.15\mu$m) protects the antenna from the harsh high altitude volcanic environment. A corrugated feedhorn with a parabolic illumination grading with a $-10.5$~dB edge taper is used to achieve low sidelobe levels. The feedhorn antenna system is simulated and designed with the mode-matching technique. The results are verified in a far-field beam pattern measurement. For the observing frequency around 94~GHz, the first sidelobe is around -20~dB with a main lobe FWHM of about 11 arcmin. With the goal of sending more stray-light to the sky, legs with a triangular roof shape are added to the secondary mirror support structure. Despite this attempt, a weak remaining feature in the beam map at the level of $\sim 1$~dB around the secondary sidelobe is likely to be attributed to the secondary mirror support structure. Measuring the weak CMB signals ($\sim$ 10 $\mu$K) poses a challenge for a close-packed array due to inter-antenna coupling. A CFRP shielding baffle is therefore added, which extends to a height of $\sim 360$~mm above the secondary mirror. Insertion and return loss measurements show that CFRP is indeed an ideal lightweight shielding material. Without loss in the antenna forward gain, the antenna cross-talk on the shortest separation of 1.4~m is measured to be $\sim -135$~dB or less. An overall antenna efficiency of about 60\% is estimated from a series of efficiency factors. The dominating loss results from the feed spill-over (efficiency $\approx 0.78$), followed by the illumination efficiency ($\sim 0.90$) and the secondary mirror blockage efficiency ($\sim 0.92$). In summary, based on the calculated and measured properties, the presented CFRP antenna is a lightweight, low side-lobe level and low noise antenna. Thus, it is appropriate for the targeted astronomical observations (Cosmic Microwave Background and galaxy cluster observations) in close-packed antenna array configurations. Currently, a 13-element compact array is used in daily routine observations.

1012.3899

1012

1012.4477_arXiv.txt

{The analysis and interpretation of the H$_2$ line emission from planetary nebulae have been done in the literature by assuming that the molecule survives only in regions where the hydrogen is neutral, as in photodissociation, neutral clumps, or shocked regions. However, there is strong observational and theoretical evidence that at least part of the H$_2$ emission is produced inside the ionized region of these objects.} {The aim of the present work is to calculate and analyze the infrared line emission of H$_2$ produced inside the ionized region of planetary nebulae using a one-dimensional photoionization code.} {The photoionization code Aangaba was improved in order to calculate the statistical population of the H$_2$ energy levels, as well as the intensity of the H$_2$ infrared emission lines in the physical conditions typical of planetary nebulae. A grid of models was obtained and the results then analyzed and compared with the observational data.} {We show that the contribution of the ionized region to the H$_2$ line emission can be important, particularly in the case of nebulae with high-temperature central stars. This result explains why H$_2$ emission is more frequently observed in bipolar planetary nebulae (Gatley's rule), since this kind of object typically has hotter stars. Collisional excitation plays an important role in populating the rovibrational levels of the electronic ground state of H$_2$ molecules. Radiative mechanisms are also important, particularly for the upper vibrational levels. Formation pumping can have minor effects on the line intensities produced by de-excitation from very high rotational levels, especially in dense and dusty environments. We included the effect of the H$_2$ molecule on the thermal equilibrium of the gas, concluding that, in the ionized region, H$_2$ only contributes to the thermal equilibrium in the case of a very high temperature of the central star or a high dust-to-gas ratio, mainly through collisional de-excitation.} {}

\defcitealias{Aleman_Gruenwald_2004}{Paper~I} Since the first detection of H$_2$ in a planetary nebula (PN) by \citet{Treffers_etal_1976}, this molecule has been detected in many PNe both in ultraviolet (UV) absorption and infrared (IR) emission \citep[e.g.][]{Beckwith_etal_1978, Zuckerman_Gatley_1988, Aspin_etal_1993, Kastner_etal_1996, Hora_etal_1999, McCandliss_etal_2007, Sterling_etal_2005, Herald_Bianchi_2004}. The H$_2$ molecules can be excited by UV photons, collisions with the gas, or by formation on excited levels. The contribution of each mechanism to the population of the H$_2$ levels depends on the physical conditions of the gas. Up to now, analyses of the excitation mechanism of the H$_2$ molecules producing the observed IR lines have been inconclusive \citep{Dinerstein_1991, Shupe_etal_1998, Hora_etal_1999, Speck_etal_2003, Rosado_Arias_2003, Likkel_etal_2006, Matsuura_etal_2007}. In the literature, the H$_2$ IR emission lines from PNe is usually analyzed under the assumption that H$_2$ molecules only exist in regions where the hydrogen is neutral, such as photodissociation regions \citep[PDRs;][]{Tielens_1993, Natta_Hollenbach_1998, Vicini_etal_1999, Bernard-Salas_Tielens_2005}, shocked regions between the expanding envelope and the wind of the progenitor \citep{Gussie_Pritchet_1988, Natta_Hollenbach_1998}, or neutral clumps inside the ionized region \citep{Beckwith_etal_1978, Gussie_Pritchet_1988, Reay_etal_1988, Tielens_1993, Schild_1995, Speck_etal_2002}. On the other hand, there is strong observational and theoretical evidence that at least part of the H$_2$ emission is produced inside the ionized region of PNe. Some authors have noticed that the morphology of some PNe shown by images taken in the H$_2$ 1-0 S(1) infrared line is very similar to those taken in [\ion{N}{II}], [\ion{S}{II}], and [\ion{O}{I}] forbidden optical lines and in hydrogen recombination lines, which are produced in the ionized region \citep{Beckwith_etal_1978, Beckwith_etal_1980, Reay_etal_1988, Webster_etal_1988, Zuckerman_Gatley_1988, Balick_etal_1991, Schild_1995, Allen_etal_1997, Guerrero_etal_2000, Lopez_etal_2000, Arias_etal_2001, Bohigas_2001, Speck_etal_2002, Speck_etal_2003}. Excitation temperatures of approximately 1000 to 2000 K are inferred from observations of H$_2$ IR emission from some PNe \citep[][and references therein]{Hora_etal_1999, Likkel_etal_2006}. Such high excitation temperatures indicate that the molecule is excited by a strong UV radiation field, since collisions or shocks cannot explain the presence of lines from excited levels with vibrational numbers over three \citep{Hora_etal_1999, Black_vanDishoeck_1987}. Furthermore, in a previous paper \citep[hereafter \citetalias{Aleman_Gruenwald_2004}]{Aleman_Gruenwald_2004}, we calculated the density of H$_2$ inside the ionized region of PNe, showing that H$_2$ can survive in a partially ionized region with moderate temperature where neutral and ionized species coexist. Since such regions can be large when the temperature of the ionizing star is high, the ionized region can be a potential contributor to the total H$_2$ line emission in some PNe. In the present paper we study the contribution of the ionized region to the observed H$_2$ IR line emission of PNe. The H$_2$ IR line intensities are calculated with a photoionization code. The effects of the temperature and luminosity of the central star, gas density, and dust-to-gas ratio on the line emission are studied. The models are described in Sect. \ref{mod}. In particular, the formalism adopted for calculating the H$_2$ energy level population and the intensities of the H$_2$ emission lines, as well as the included H$_2$ excitation mechanisms are described in detail. Results are discussed in Sect. \ref{res}. A summary of the conclusions and final comments are presented in Sect. \ref{final}.

\label{final} In the present work we have studied the H$_2$ infrared emission from the ionized region of PNe. For this, we used the one-dimensional photoionization code Aangaba, in which we included the physics and chemistry of the H$_2$ molecule. This powerful tool can now be used to study the H$_2$ density, level population, line emission, processes of formation and destruction, and mechanisms of population and depopulation of the molecular energy levels as a function of the distance from the ionizing source of an ionized gaseous nebula. Although there is observational evidence that at least part of the H$_2$ 1-0 S(1) line emission may originate inside the ionized region of PNe, the published studies of the molecular hydrogen emission of PNe do not usually take into account the contribution of the ionized region to this emission. One of the important conclusions of this work is that the H$_2$ emission of the ionized work can contribute significantly to the total emission observed in PNe, particularly for PNe with high $T_\star$. Comparison between the calculated and observed H$_2$ 1-0 S(1)/Br$\gamma$ ratio shows that the emission of the ionized region can be responsible for a substantial fraction of the total H$_2$ emission in such cases. Therefore it is important to use a code where the neutral and the ionized region are taken into account self-consistently. In the ionized region, the H$_2$ IR emission lines are produced predominantly in the TZ, where the atomic lines [\ion{N}{II}], [\ion{O}{I}], and [\ion{S}{II}] are significantly produced. The partially ionized and warm gas in the TZ favors the formation and survival of H$_2$ molecules, as well as its IR line emission. The temperature of the central star is an important factor for the H$_2$ density and IR line intensity, since hotter stars produce more high-energy photons than colder stars. These photons can penetrate deep into the nebulae, producing the TZ. The 1-0 S(1) line intensity increases strongly with the increase in $T_{\star}$ in our models. This result agrees with the correlation between the detection of H$_2$ in PNe and their Zanstra temperatures found by \citet{Phillips_2006}. Furthermore, this can explain why the detection of this H$_2$ line is more common in bipolar PNe (Gatley's rule), given that these objects typically have higher $T_{\star}$. Although \citet{Reay_etal_1988} suggest that the correlation between the intensity of the 1-0 S(1) and [\ion{O}{I}] lines in PNe is due to the existence of clumps, such correlation can also be naturally explained by both lines being produced in the TZ. The most intense H$_2$ lines are emitted in the range 1 to 29 $\mu$m, in the bands 0-0, 1-0, 1-1, 2-0, 2-1, 3-1, 4-1, 4-2, and 5-3. Several of these intense lines have already been detected in PNe. The 1-0 S(1) line is one of the more intense lines. The fraction of the 1-0 S(1) line to the total H$_2$ IR line emission is around 1\% (within a factor of less than 10) in all our models. Both collisions and UV pumping play important roles in the excitation of H$_2$ infrared lines in the ionized region. The effect of the excitation by UV pumping is important for levels with $v > 3$. The relative importance of collisions over UV pumping increases with the increase in $T_{\star}$ and $n_\mathrm{H}$. Grain surface reaction and associative detachment may only be significant for very high levels ($v > 5$ and $J > 6$), particularly in the cases of denser nebula or colder central stars. Grain surface reaction may be important for lower levels in the presence of large amounts of dust. The effect of H$_2$ on the thermal equilibrium is insignificant in most models, except in the cases of a very high central star temperature or dust density, where the molecule cools the gas in the outer zone of the TZ, primarily through collisional excitation.

1012.4477

1012

1012.1855_arXiv.txt

We present the rest-frame optical morphologies of active galactic nucleus (AGN) host galaxies at $1.5<z<3$, using near-infrared imaging from the \textit{Hubble Space Telescope} Wide Field Camera 3, the first such study of AGN host galaxies at these redshifts. The AGN are X-ray selected from the \textit{Chandra} Deep Field South and have typical luminosities of $10^{42} < L_{\rm X} < 10^{44}$\ergs. Accreting black holes in this luminosity and redshift range account for a substantial fraction of the total space density and black hole mass growth over cosmic time; they thus represent an important mode of black hole growth in the universe. We find that the majority ($\sim80\%$) of the host galaxies of these AGN have low \sersic\ indices indicative of disk-dominated light profiles, suggesting that secular processes govern a significant fraction of the cosmic growth of black holes. That is, many black holes in the present-day universe grew much of their mass in disk-dominated galaxies and not in early-type galaxies or major mergers. The properties of the AGN host galaxies are furthermore indistinguishable from their parent galaxy population and we find no strong evolution in either effective radii or morphological mix between $z\sim2$ and $z\sim0.05$.

\label{sec:intro} Host galaxy morphology is a key parameter in the joint formation of galaxies and supermassive black holes via active black hole growth phases \citep{2010ApJ...711..284S}. In the local universe, black hole growth is associated with very different evolutionary stages of galaxies: early-type host galaxies feature black hole accretion in a specific time window after a rapidly suppressed burst of star formation that was induced by a merger event \citep{2007MNRAS.382.1415S, 2010ApJ...714L.108S, 2010ApJ...711..284S, 2007MNRAS.381..543W, 2008ApJ...673..715C}, while black hole growth in late-type galaxies, which dominates active galactic nucleus (AGN) host galaxy sample by number, are massive, stable disk galaxies with no obvious recent perturbations to star formation \citep{2010ApJ...711..284S}. That is major mergers do not seem to make up a large part of the AGN host galaxy population at low redshift, although puzzlingly the fraction of AGN exhibiting signs of recent or ongoing interactions increases at the highest luminosities in unbiased, hard X-ray selected samples such as the \Swift\ BAT sample \citep{2009ApJ...692L..19S, 2010ApJ...716L.125K}. In any case, accretion at $z\sim0$ represent a negligible fraction of cosmic black hole growth, most of which occurred at high redshift. \begin{figure*} \begin{center} \includegraphics[angle=90,width=0.163\textwidth]{fig1a.ps} \includegraphics[angle=90,width=0.163\textwidth]{fig1b.ps} \includegraphics[angle=90,width=0.163\textwidth]{fig1c.ps}\\ \includegraphics[angle=90,width=0.163\textwidth]{fig1d.ps} \includegraphics[angle=90,width=0.163\textwidth]{fig1e.ps} \includegraphics[angle=90,width=0.163\textwidth]{fig1f.ps} \caption{Sample $H$-band (F160W) cutouts measuring $5.6\arcsec \times 5.6\arcsec$ of the X-ray selected AGN host galaxies in our sample. In the top-left of each cutout, we give the corresponding ID number (see Table \ref{tab:agn}). In the top row are galaxy-dominated AGN host galaxies while the left and middle of the bottom row are AGN-dominated sources. In the bottom-right panel, we show the empirical PSF generated from stars in the field. \label{fig:gallery}} \end{center} \end{figure*} The properties of AGN host galaxies during the peak epoch of both star formation and black hole growth at $z\sim2$ have until very recently remained virtually inaccessible. Ground-based imaging does not offer sufficient angular resolution to resolve the AGN host galaxies, while most \textit{Hubble Space Telescope} imaging was in optical bands that translates to the rest-frame ultraviolet at $z\sim2$. With the installation of the new Wide Field Camera 3 (WFC3) on the \textit{Hubble Space Telescope}, we now have available ultra-deep, high resolution near-infrared images of AGN at $z\sim2$ which allow us to study the \textit{rest-frame optical} properties of their host galaxies in detail. The $F160W$ ($H$-band) filter corresponds approximately to the $V$-band at $z\sim2$, with spatial resolution comparable to or better than Sloan Digital Sky imaging at $z\sim0.05$\footnote{For SDSS, the $1.^{''}4$ median seeing at $z=0.05$ corresponds to 1.36 kpc, while for \textit{HST} WFC3/IR, the $0.^{''}13$ (undersampled) pixel scale at $z=2$ corresponds to 1.09 kpc.}. In this \textit{Letter}, we present the restframe optical morphologies of moderate luminosity AGN ($10^{42} < L_{\rm X} < 10^{44}$\ergs; corresponding to $-23 \lesssim M_{\rm V} \lesssim -18$) during the peak epoch of growth at $z\sim2$. AGN with these luminosities represent a significant fraction of the cosmic black hole growth in terms of both number density and in X-ray light emitted \citep[e.g.,][]{2003ApJ...598..886U, 2005A&A...441..417H}. Throughout this \textit{Letter}, we assume a $\Lambda$CDM cosmology with $h_{0}=0.7$, $\Omega_{m}=0.27$ and $\Omega_{\Lambda}=0.73$, in agreement with the most recent cosmological observations \citep{2009ApJS..180..225H}. \begin{figure} \begin{center} \includegraphics[angle=90, width=0.49\textwidth]{fig2b.ps}\\ \includegraphics[angle=90, width=0.49\textwidth]{fig2d.ps} \caption{An example of an F160W images of AGN host galaxies and the \galfit\ output for a galaxy dominated by starlight but containing a nuclear point source. Below the image are the three best-fit models (\sersic\ only, \sersic+PSF and PSF only), with a logarithmic stretch, and below that are the corresponding residual images. The PSF-only residuals clearly show a resolved component that the PSF-only fit could not accommodate. \label{fig:examples}} \end{center} \end{figure} \begin{deluxetable*}{lcccccccc} \tablecolumns{9} \tablewidth{0pc} \tabletypesize{\scriptsize} \tablecaption{X-ray-selected AGN host galaxies in the WFC3/IR ERS field} \tablehead{ \colhead{ID$^1$} & \colhead{RA} & \colhead{Dec} & \colhead{Redshift} & \colhead{Type$^2$} & \colhead{$log_{10}(L_{\rm{X, obs}})$} & \colhead{$log_{10}(M_{\rm{stellar}})$$^3$} & \colhead{Eddington} & \colhead{$M_{\rm V}$} \\ \colhead{} & \colhead{(J2000)} & \colhead{(J2000)} & \colhead{} & \colhead{} & \colhead{\ergs} & \colhead{\Msun} & \colhead{Ratio$^4$} & \colhead{AB mag} } \startdata 49190 & 03 32 19.9 &-27 45 17.9 & 2.424 & phot & 44.05 & 10.24 & 0.84 & -22.17 \\ 50057 & 03 32 06.7 &-27 44 55.1 & 2.296 & phot & 42.94 & 10.25 & 0.06 & -22.05\\ 50333 & 03 32 03.0 &-27 44 50.0 & 2.573 & spec & 44.06 & 10.39 & 0.57 &-22.79\\ 50634 & 03 32 04.0 &-27 44 41.5 & 3.618 & phot & 43.52 & 10.23 & 0.25 &-22.85\\ 52141 & 03 32 10.9 &-27 44 14.9 & 1.613 & spec & 44.35 & 10.08 & 2.60 &-22.97\\ 52399 & 03 32 14.8 &-27 44 02.6 & 1.527 & phot & 42.78 & 10.32 & 0.04 &-19.95\\ 53849 & 03 32 15.1 &-27 43 35.3 & 1.691 & phot & 43.02 & 10.81 & 0.02 & -23.16\\ 54369 & 03 32 01.6 &-27 43 27.0 & 2.720 & spec & 44.47 & 10.41 & 1.38 & -22.23\\ 55062 & 03 32 25.7 &-27 43 05.7 & 2.291 & spec & 44.39 & 10.89 & 0.30 & -22.85\\ 55620 & 03 32 24.2 &-27 42 57.7 & 2.303 & spec & 43.39 & 10.54 & 0.08 & -22.57\\ 56112 & 03 32 20.0 &-27 42 43.6 & 2.733 & phot & 43.86 & 10.07 & 0.87 & -21.35\\ 56769 & 03 32 20.2 &-27 42 27.2 & 2.773 & phot & 43.58 & \nodata & \nodata & -21.06\\ 56954 & 03 32 14.1 &-27 42 30.1 & 2.026 & spec & 42.99 & 10.49 & 0.04 & -22.07\\ 57420 & 03 32 25.2 &-27 42 18.8 & 1.617 & spec & 44.00 & 11.05 & 0.08 & -23.72\\ 57805 & 03 32 15.0 &-27 42 25.0 & 1.895 & phot & 43.80 & 11.48 & 0.02 & -25.13\\ 57859 & 03 32 15.8 &-27 42 07.6 & 2.779 & phot & 43.32 & 10.71 & 0.04 & -22.94\\ 58039 & 03 32 33.9 &-27 42 04.1 & 1.936 & phot & 42.85 & 10.85 & 0.01 &-23.57\\ 58224 & 03 32 15.2 &-27 41 58.6 & 2.402 & spec & 43.47 & 10.56 & 0.09 & -22.67\\ 58330 & 03 32 05.0 &-27 42 02.7 & 2.062 & phot & 43.16 & 10.40 & 0.07 &-23.29\\ 58509 & 03 32 12.7 &-27 41 49.0 & 1.490 & phot & 43.02 & 9.91 & 0.19 & -20.23\\ 58657 & 03 32 08.3 &-27 41 53.5 & 2.470 & spec & 43.68 & 10.35 & 0.26 &-21.68\\ 59060 & 03 32 17.1 &-27 41 37.0 & 2.193 & phot & 43.85 & 10.57 & 0.21 & -21.79\\ 63732 & 03 32 35.4 &-27 40 02.7 & 1.490 & phot & 42.60 & 10.27 & 0.03 &-22.13\\ \enddata \label{tab:agn} \tablenotetext{1}{Catalog ID, see \cite{2010ApJS..189..270C}.} \tablenotetext{2}{Type of redshift: phot - photometric redshift, spec - spectroscopic redshift; see \cite{2010ApJS..189..270C}.} \tablenotetext{3}{Stellar mass as computed by \cite{2010ApJ...721L..38C}.} \tablenotetext{4}{Black hole mass is calculated from this stellar mass following the relation of \cite{2004ApJ...604L..89H}. Eddington ratios assume an X-ray to bolometric correction factor of 20. Since a substantial fraction of the total stellar mass is likely in a disk, the black hole masses are overestimates, which means the Eddington ratios are lower limits. } \end{deluxetable*}

\label{sec:discussion} We have obtained the first clear view of the rest-frame optical morphologies of AGN host galaxies with Seyfert-like luminosities ($10^{42} < L_{\rm X} < 10^{44}$\ergs) at $z\sim2$ using the new \textit{Hubble Space Telescope} WFC3/IR in the ERS portion of the GOODS-S field. Fits to the host galaxy surface brightness profiles reveals that: \begin{enumerate} \item The majority of these AGN host galaxies have low \sersic\ indices, implying the bulk of the host galaxy light comes from a disk. \item The host galaxy structural parameters (\sersic\ index and effective radius) do not appear to be significantly different from a comparison sample of inactive galaxies matched in redshift and luminosity. \item The distribution of \sersic\ indices implies that high redshift AGN host galaxies have very similar morphologies to local AGN host galaxies, \textit{i.e.}, few early types but many late types. The typical effective radii are also similar to those of local AGN host galaxies. \end{enumerate} These AGN host galaxies are a significant fraction of the total AGN population by number density and in terms of light emitted by accretion \citep[][]{2003ApJ...598..886U, 2005A&A...441..417H}. Using the X-ray luminosity function of \cite{2003ApJ...598..886U} and evolution, obscuration distribution and bolometric correction as described by \cite{2009ApJ...696..110T}, we estimate the black hole growth in this population in the $z=1.5-3$ range spanned by our sample represents 10--17\% of the total black hole growth over cosmic history. Excluding the most massive black holes, which get most of their mass in quasar-luminosity events triggered by mergers \citep{2010Sci...328..600T}, 23--40\% of black hole growth occurs in a secular mode driven by internal processes in the host galaxy. Since disks also dominate the AGN host galaxy population at $z\sim0$ \citep{2010ApJ...711..284S}, where quasar-mode growth is unimportant, an even larger fraction of \textit{all} black hole growth over cosmic history appears to take place in disk galaxies. The results presented here show that moderate luminosity AGN host galaxies at $z\sim2$ and $z\sim0$ are remarkably similar. The high fraction of AGN host galaxies with disk-like light profiles is difficult to reconcile with the expectation of black holes growing jointly with stellar bulges during special phases of their evolution, such as major mergers envisioned in many simulations \citep[e.g.,][]{2005MNRAS.361..776S, 2005ApJ...630..705H, 2008ApJS..175..356H}. The disk morphologies of the host galaxies point instead to secular processes \citep[e.g.,][]{2004ARA&A..42..603K} as most common growth mode. The fact that AGN host galaxies are indistiguishable from the $z\sim2$ comparison sample in terms of their \sersic\ indices and effective radii further supports the role of secular growth. This is very different from the high-luminosity (quasar) population at the same redshift, which does seem to be driven by major mergers \citep{2010Sci...328..600T}. Thus in the high redshift universe, there appear to be two distinctly different modes of black hole growth for high- and low-luminosity AGN. The fact that the majority of black hole growth in this population --- and by extension a significant fraction of cosmic black hole growth --- occurs in a galaxy substantial disk means that it is not associated with major mergers. This raises interesting questions regarding the origin and relevance of the relationship between galaxy and black hole mass \citep{2000ApJ...539L..13G, 2000ApJ...539L...9F, 2007ApJ...671.1098P, 2010arXiv1006.0482J}. This secular black hole growth must still be self-regulated in some way that preserves the correlation between black hole mass and bulge mass.

1012.1855

1012

1012.1199.txt

We study, by numerical methods, the time evolution of scalar perturbations in radiation era of Randall-Sundrum braneworld cosmology. Our results confirm an existence of the enhancement of perturbation amplitudes (near horizon crossing), discovered recently. We suggest the approximate solution of equations of the perturbation theory in the high energy regime, which predicts that the enhancement factor is asymptotically constant, as a function of scale. We discuss the application of this result for the problem of primordial black hole production in braneworld cosmology.

During last decade, braneworld cosmological scenarios, in which our 4D Universe is realized as a hypersurface embedded in a higher-dimensional spacetime have attracted much attention. In first scenarios of this kind, suggested as early as in 1980's \cite{Akama:1982jy, Rubakov:1983bb}, it had been shown that matter fields can be confined to a field-theoretical domain wall (topological defect) in a world with non-compact extra dimensions. The progress in string theory in subsequent years, especially the discovery of D-branes, has revived interest to the idea of braneworlds. In general, the string theory is quite promising, it may provide an unified description of gauge interactions and gravity. In the present context, it is most important that it predicts the existence of $p$-branes, ($p+1$)-dimensional sub-manifolds of the 10 (or 11) - dimensional spacetime on which open strings end. Gauge particles and fermions which correspond to string end points can only move along these $p$-branes, while gravitons can propagate in the full spacetime (``bulk''). It is tempting to assume that our ($3+1$)-dimensional spacetime is such a 3-brane. If only gravity can probe the bulk, the extra dimensions can be very large (in comparison with the smallest length scale tested, so far, in particle physics, $\sim 10^{-16}\;$cm). It had been assumed in \cite{ArkaniHamed:1998nn}, that the extra dimensions are compact, in analogy with the old Kaluza-Klein (KK) picture \cite{KK}. Slightly later, in works by Randall and Sundrum \cite{Randall:1999ee, Randall:1999vf}, it was pointed out that this condition is not necessary and the extra dimension may be even non-compact. The Randall-Sundrum (RS) model is of particular interest due to its relative simplicity, in spite of the fact that it includes nontrivial gravitational dynamics. In the RS2 model \cite{Randall:1999vf} a single brane is embedded in a anti - de Sitter (AdS) bulk and, although the 5th dimension extends infinitely, the warped structure of the bulk geometry (i.e., the curvature of the bulk spacetime) leads to a recovery of the standard General Relativity (GR) on the brane at scales larger than the bulk curvature scale $\ell$. In particular, Newton's law is recovered at large distances and the Friedmann's equation for the evolution of the Universe is obtained at low energy. At high energies, i.e., in the very early Universe, the Friedmann equation differs substantially from GR by a correction term which is proportional to $\rho/\sigma$, where $\rho$ is the density of brane matter, and $\sigma$ is the brane tension. This term leads to a faster Hubble expansion at high energies. Inflationary expansion of the Universe is also modified in brane cosmology: the evolution of the inflaton field is more strongly damped, and the brane Universe inflates at much faster rate than what is expected from standard cosmology. Another important effect at high energies is the excitation of KK-modes which escape from our brane into the 5D bulk, leading, in particular, to the suppression of the power spectrum of inflationary gravitational wave background. Cosmological perturbation theory in braneworld cosmology also has some distinct features \cite{Mukohyama:2000ui, Kodama:2000fa, vandeBruck:2000ju, Langlois:2000ph, Koyama:2000cc, Deffayet:2002fn}. The equations of the perturbation theory contain high-energy corrections ($\sim\rho/\sigma$) similar to those in the Friedmann equation and, in addition, the correction terms arising from the fluctuations of the bulk geometry. Perturbations on brane, e.g., the scalar perturbations (which we are interested in) are coupled with the bulk perturbations. Technically, in a case of the scalar perturbations and AdS bulk, the problem is reduced to the solution of a system of equations for the density contrast variable and the so-called master variable (it appears that all quantities describing the bulk perturbations are written in terms of this variable \cite{Mukohyama:2000ui, Mukohyama:2001yp, Kodama:2000fa}). In the context of braneworld models, a question about existence and evolution laws of the higher-dimensional black holes is very interesting and important. In a model with the 5th large extra dimension, a physically meaningful black hole solution is the 5D-Schwarzschild \cite{Tangherlini:1963bw, Myers:1986un}, if the horizon size is sufficiently small compared with an effective size of the extra dimension. Really, it is natural to assume that primordial braneworld black holes formed in the early Universe with a horizon size $r_s \ll \ell$ would be described by a 5D Schwarzschild metric because in this case the AdS curvature has very little effect on the geometry. Numerical calculations support the existence of static solutions for such small $r_s$ \cite{Kudoh:2003xz}. However, the results of these calculations cannot be extrapolated to the case $r_s\sim \ell$. Unfortunately, an exact solution representing a localized and stable black hole is known only in 4D braneworld model \cite{Emparan:1999wa}, whereas the corresponding solution in the 5D braneworld model has not been found. The process of the gravitational collapse on the brane is very complicate, due to, in particular, gravitational interaction between the brane and the bulk (see, e.g., \cite{Maartens:2010ar}). Even in the simplest case of RS-type brane, and Oppenheimer-Snyder (OS) - like collapse, braneworld gravity introduces important new features in the black hole formation process (the high energy- and KK-corrections to the field equations of GR, i.e., the same corrections which affect the expansion of the early Universe, are also efficient here). These features lead to a non-static exterior of the black hole \cite{Bruni:2001fd} in the case of the OS-collapse. Moreover, there are arguments \cite{Tanaka:2002rb, Emparan:2002px} based on AdS/CFT-correspondence, that the non-static behavior exists also in a general collapse. If, really, the black hole solutions in braneworld scenarios, for a black hole larger than AdS radius, are quite different from those in 4D GR (i.e., if, as authors of \cite{Tanaka:2002rb, Emparan:2002px} argue, these solutions are necessarily non-static and predict short lifetime of large black holes due to the strongly enhanced evaporation), there is an unique chance to probe the extra dimension by astronomical observations of massive black holes. Predictions for an evolution of the small ($r_s\ll \ell)$ black holes are less dramatic (and less speculative). The differences from the 4D case are reduced to a larger probability of accretion, in the high energy regime (due to the fact that in this regime the radiation density is proportional to $t^{-1}$ rather than $t^{-2}$) and to a relative increase of the primordial black hole (PBH) lifetime, for a given initial mass. In particular, initial mass of PBHs evaporating today can be $10^9 - 10^{10}\;$g rather than $\approx 10^{15}\;$g as predicted by GR. For the PBHs having small masses, there are astrophysical constraints on their abundance, based, e.g., on studies of extragalactic photon and neutrino backgrounds. These constraints give, as usual, the information about primordial density perturbations (we assume that PBHs form from these perturbations). For an extraction of this information one must know the evolution of these perturbations in radiation era. In the recent work by Cardoso {\it et al} \cite{Cardoso:2007zh} it had been shown that the density perturbations with short wavelengths are amplified during horizon re-entry. The magnitude of this enhancement depends, clearly, on a scale of the density perturbations. The smaller is the scale, the earlier the perturbation crosses horizon, and, if comoving wave number $k$ is larger than some critical value $k_c$, this crossing happens at high energy regime. The straightforward calculation of the enhancement factor, in the region of scales which are relevant for PBHs with small masses, evaporating near today, is quite difficult, even numerically, due to a very complicate machinery of cosmological perturbation theory in braneworld cosmology. In the present paper we study the dependence of the enhancement factor on the comoving size of the density perturbations. We carried out detailed numerical calculations of gauge invariant amplitudes of curvature perturbations as functions of the scale factor and the corresponding enhancement factors. We found the approximate solution (of the equations for perturbation amplitudes), describing the time evolution of the amplitudes near horizon crossing. According to this solution, the magnitude of the enhancement factor doesn't depend, in the high energy region, on the comoving scale. Using this conclusion it is possible to calculate the enhanced perturbation amplitudes for arbitrarily small scale. The plan of the paper is as follows. In the second section the equations of perturbation theory in RS2 braneworld cosmology which are necessary for curvature perturbation calculations are given. In the third section, the approximate solution of these equations in the high energy limit is suggested. In the fourth section the main relations characterizing the PBH evolution in the RS2 braneworld are briefly reviewed. The scheme used in the numerical calculations is presented in Sec. \ref{sec-scheme}. The results of the paper and conclusions are summarized in the last section.

\label{sec-results} The main results of the numerical calculations are shown in Figs. \ref{fig-Om-3kc} - \ref{fig-enh}. Figs. \ref{fig-Om-3kc}, \ref{fig-3kc-delt} show, for a given value of $k$, $k=3k_c$, the evolution, near horizon crossing, of three variables: $\Omega_b$ (Fig. \ref{fig-Om-3kc}), $\Delta$ (Fig. \ref{fig-3kc-delt}, upper panel) and $\zeta$ (Fig. \ref{fig-3kc-delt}, lower panel). It is clearly seen that all three variables rise with $a$, up to $a\sim 3 a_*$. It is also seen from Fig. \ref{fig-3kc-delt} that the corresponding rise of $\Delta$ and $\zeta$ in GR is weaker, and, as a result, we have an enhancement. It follows also, from Fig. \ref{fig-3kc-delt}, that the enhancement is not zero in the approximation $\Omega_b=0$, when there are no KK corrections in the equations of the 5D perturbation theory. In Figs. \ref{fig-zeta-k}, \ref{fig-Omega} it is shown how the evolution curves for $\zeta$ and $\Omega_b$ change with an increase of $k$. It is seen, from Fig. \ref{fig-zeta-k}, that the enhancement grows with $k$, but there is a clear tendency of a slowdown of this growth at $k > 10 k_c$. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \begin{figure} \center % \includegraphics[width=8 cm]{zeta-diff-k.eps} \caption{ \label{fig-zeta-k} The result of the numerical calculation of $\zeta$ for different values of $k$.} % \end{figure} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \begin{figure} \center % \includegraphics[width=8 cm]{Omega.eps} \caption{ \label{fig-Omega} The result of the numerical calculation of $\Omega_b$ for different values of $k$, normalized to $\zeta_k=1$ in the super-horizon regime. Arrows show the horizon crossing time ($k=aH$) for each mode.} % \end{figure} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% Following \cite{Cardoso:2007zh} we define the factors that show the degree of enhancement of the perturbation amplitudes: \begin{equation} {\cal Q}_{eff} = \frac {\Delta_{eff}} {\Delta_{GR}} \;, \; {\cal Q}_{\mathcal{E}} = \frac {\Delta_{5D}} {\Delta_{eff}} \;, \; {\cal Q}_{5D} = \frac {\Delta_{5D}} {\Delta_{GR}} = {\cal Q}_{eff} {\cal Q}_{\mathcal{E}}. \end{equation} In a case of the effective theory ($\Omega_b=0$), the enhancement reaches an asymptotic value, ${\cal Q}_{eff} \approx 3$, at $k\sim 100 k_c$. However, the direct calculation of ${\cal Q}_{5D}$ (or, equivalently, ${\cal Q}_{\mathcal{E}}$) for very large wave numbers $k\gg k_c$ is not easy, due to a quite complicate behavior of $\Omega$ in the bulk (see Fig. \ref{fig-grid} for an illustration: the larger value of $k$, the more frequent are the oscillations in the bulk). Due to limitations of computing resources, we have been able to make direct calculations in 5D case only for a limited range of $k\lesssim 30 k_c$. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \begin{figure}[!b] \center % \includegraphics[width=7.5 cm]{enh.eps} \caption{ \label{fig-enh} Enhancement factors that show the degree of increasing of the perturbation amplitude after horizon entry. From bottom to top, curves show the enhancement of the amplitude of 5-dimensional calculation compared to the effective one, effective theory compared to General Relativity result and 5-dimensional calculation compared to General Relativity result.} % \end{figure} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% In order to study the PBH production for masses $M_{BH}^*(t_0) \sim 10^9\;$g (such PBHs, as we have seen in Sec. \ref{sec-5D}, evaporate near today, if the value of $\ell$ is close to its upper bound (\ref{ell-bound})), we need information about cosmological perturbations for $k \gtrsim 10^6 k_c$. To perform calculations for such large wave numbers, we have used an approximate approach according to which $\Omega_b(a/a_*)$ has the same form (and is given by Eq. (\ref{ff})) for all large wave numbers (see the discussion in Sec. \ref{Sec-f}). In these calculations we have used, for the required function $f_\Omega(a/a_*)$, the corresponding function obtained from the direct numerical calculation for $k=30 k_c$. Using this approach, we have calculated the enhancement factors for large values of $k$ ($k \gtrsim 30 k_c$). The results of the calculation are shown in Fig. \ref{fig-enh}. In summary, we stressed in this paper that, in RS2 brane cosmology, the PBHs of relatively small mass (the concentration of which in space can be constrained by cosmological arguments) form in the high energy regime, and the corresponding comoving wave numbers are very large, $k \sim (10^6 - 10^7) k_c$. We thoroughly studied, by numerical methods, the evolution of scalar perturbation amplitudes (those needed for calculations of the PBH production) near horizon crossing, for a wide range of comoving scales. We confirmed the main conclusion of \cite{Cardoso:2007zh} according to which amplitudes of the curvature perturbation get enhanced after horizon re-entry (before a beginning of the oscillation phase). We developed an approximate phenomenological approach for calculations of the perturbation amplitudes for very small scales, where the direct numerical methods are powerless. We argued, using this approach, that in the asymptotic limit of high energies, the enhancement factor is constant as a function of the perturbation scale. We presented details of the numerical scheme (based on the pseudo-spectral method) which is used for a treating of scalar cosmological perturbations on the brane and in the bulk. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%

1012.1199

1012

1012.4641_arXiv.txt

{The post-processing of passively advected Lagrangian tracer particles \cite{nagataki1997a} is still the most common way for obtaining detailed nucleosynthetic yield predictions of Type Ia supernova (SN Ia) hydrodynamical simulations. Historically, tracer particles of constant mass are employed. However, intermediate mass elements, such as e.g. Ne, Mg, Al, or Si, are typically synthesized in the outer layers of SNe Ia, where due to the lower initial density a constant mass tracer distribution results in poor resolution of the spatial morphology of the abundance distribution. We show how to alleviate this problem with a suitably chosen distribution of variable tracer particle masses. We also present results of the convergence of integrated nucleosynthetic yields with increasing tracer particle number. We find that the yields of the most abundant species (mass fraction > $10^{-5}$) are reasonably well predicted for a tracer number as small as 32 per axis and direction. Convergence for isotopes produced in regions where a constant tracer mass implementation results in poor spatial resolution can be improved by suitably choosing tracers of variable mass.} \FullConference{11th Symposium on Nuclei in the Cosmos\\ 19-23 July 2010 \\ Heidelberg, Germany.} \begin{document}

Type Ia supernovae (SN Ia) are believed to be thermonuclear explosions of white dwarf stars. The relative abundances of nuclei synthesized in SN Ia explosions are dependent on the explosion model and play an important role in understanding the chemical evolution of our Galaxy (e.g. \cite{matteucci2009a}). Multi-dimensional numerical simulations of SN Ia explosions have been carried out by several groups for a range of different explosion models (see e.g. \cite{roepke2007c,roepke2007b,bravo2008a,jordan2008a,meakin2009a,pakmor2010a,fink2010a,sim2010a}). Synthetic light curves and spectra not only depend sensitively on the amount and location of the radioactive nuclei that decay and reheat the ejecta (such as e.g. \nuc{56}{Ni} or \nuc{57}{Ni}) but also on the amount and location of many other elements such as e.g. Mg, Ca, Ti or Cr, whose (partially ionized) atoms interact with the radiation. For any explosion model, detailed predictions of the isotopic composition of the ejecta are therefore needed. This is commonly done with the tracer particle method (e.g. \cite{travaglio2004a,brown2005a,roepke2006a,fink2010a,maeda2010a}), for which a large number of tracer particles is placed into the star. The particles are carried along by the flow, recording the local thermodynamic conditions as a function of time. These ``trajectories'' are then post-processed with a nuclear reaction network and the resulting yields are assigned to the tracer particles final position and velocity, weighted by the mass the particle represents. Historically, tracer particles of equal (constant) mass are used throughout.

The tracer particle method yields appear to begin to converge for tracer particle numbers greater than $\sim$32 per axis and direction (see Fig.~\ref{fig3}). The yield convergence for nucleosynthetic products of incomplete burning, such as e.g. Mg or Al, can be improved upon by using tracer particles of variable mass. These variable mass tracers can greatly improve the spatial resolution of the lower density nucleosynthetic yield distribution, which could lower the number of particles required to obtain converged light curves and spectra (see e.g. \cite{kromer2009a}). For further details, please see the journal article associated with this poster \cite{seitenzahl2010a}.

1012.4641

1012

1012.0254_arXiv.txt

Planetary migration is essential to explain the observed mass-period relation for exoplanets. Without some stopping mechanism, the tidal, resonant interaction between planets and their gaseous disc generally causes the planets to migrate inward so efficiently that they plunge into the host star within the gaseous disc lifetime ($\sim $ 1-3 Myrs). We investigate planetary migration by analytically calculating the migration rate and time within self-consistently computed, radiatively heated discs around M stars in which the effects of dust settling are included. We show that dust settling lowers the disc temperature and raises the gas density in the mid-plane. This inescapable evolution of disc structure speeds up type I planetary migration for lower mass bodies by up to a factor of about 2. We also examine the effects of dust settling on the gap-opening mass and type II migration, and find that the gap-opening mass is reduced by a factor of 2 and type II migration becomes slower by a factor of 2. While dust settling can somewhat alleviate the problem of planetary migration for more massive planets, the more rapid migration of low mass planets and planetary cores requires a robust slowing mechanism.

Planet formation involves the interaction between protoplanets and gaseous protoplanetary discs. This interaction drives radial orbital drifts for protoplanets \citep[hereafter GT80]{gt79,gt80}, possibly resulting in the diversity of exoplanetary orbital properties captured in the observed mass-period relation \citep{us07}. In standard disc models, planets lose their net angular momentum by resonant interaction with their natal discs and migrate inward. This is confirmed by analytical calculations \citep[hereafter, TTW02]{ward97,ttw02} and numerical simulations \citep{npmk00,kdh01,dhk02,dkh03,blom03} in 2D and 3D discs. It is well established that the driving force of planetary migration is exerted only at Lindblad and corotation resonances (GT80). While Lindblad torques excite and propagate density waves in discs \citep{ward97}, corotation torques do not. Due to the accumulation of angular momentum around the corotation region in discs, without any removal mechanisms such as disc viscosity \citep{ward91,m01,m02}, the corotation torque saturates (i.e. becomes ineffective). The problem of migration in the standard disc models is that, without some sort of slowing mechanism, this tidal interaction between a planet and the gaseous disc is so efficient that the migration timescale is much shorter than the disc lifetime ($\sim $ 1 - 3 Myrs). This fast mode of migration, known as type I migration, is applicable for low mass planets which cannot open up a gap in discs. Based on the core accretion scenario, which is the most accepted theory for planet formation, cores of gas giants and rocky planets are classified as type I migrators. Massive planets such as gas giants can open a gap and undergo the so-called type II migration, which is about two orders of magnitude slower than type I migration \citep[e.g.][]{ward97}. Population synthesis models confirm that type I migration should be about two orders of magnitude slower than expected in order to qualitatively reproduce the observed mass-period relation \citep{il08v,mab09}. One of the greatest challenges in planetary formation and migration is to identify what physical process(es) succeeds in reproducing the diversity of exoplanets and under what conditions planetary systems similar to our Solar System are formed. This long-standing problem can be resolved by paying close attention to the properties of discs. Tidal torque strongly depends on the distributions of gas and temperature in discs and is sensitive to inhomogeneities in these quantities \citep[hereafter, HP10, and references herein]{hp10}. Two sources of disc inhomogeneity have been proposed in the literature: dead zones and ice lines. Dead zones, regions with high density and low turbulence \citep[i.e. where the magneto-rotational instability (MRI) is quenched;][]{g96,mp06}, drastically affect these distributions and eventually provide two of the most robust slowing mechanisms \citep[HP10]{mpt07}. Both of them are the consequence of the difference in turbulence between the active and dead zones. One barrier arises by piling up gas at a dead zone's outer boundary which is a consequence of the time-dependent, viscous evolution of discs that have strong variations in the turbulent intensity \citep[$\alpha$;][]{mpt07,mpt09}. The other is the result of a positive temperature gradient (HP10). This radial temperature gradient in the dead zone is produced by the back heating from a dusty wall which is left in the active region due to the enhanced dust settling in the dead zone \citep[hereafter, Paper I]{hp09b}. The variation in turbulence with radial direction is the focus in the above treatments. The vertical variation in turbulence, resulting in layered structures: the MRI-active surface and the MRI-dead regions \citep{g96} is also important. Combined with the layered structures, the ice lines act as a barrier for type I migration \citep[e.g.][]{il08v}. This barrier arises because at the ice lines, the ice-coated dust density suddenly increases due to the low disc temperature. Consequently, the electron number density in the surface drops sharply due to the absorption by such dust. Since electrons are argents for gas to be MRI-active, the gas density of the MRI-active surface region falls and gas piles up in the MRI-dead region. As a result, a radial positive gradient for the surface density appears, which acts as a barrier for type I migration \citep{mmcf06,pp09a}. In this paper, we step back and investigate the general effects of dust settling on planetary migration in homogeneous discs. Although dust settling is observationally confirmed to be ubiquitous in discs around young stars (Paper I, references herein), there is no comprehensive study so far on the effect that this has on planetary migration. Furthermore, dust settling is one of the important processes for planet formation \citep[e.g.][]{bdb10}. In order to proceed, we compute detailed, radiatively heated disc models around an M star as described in Paper I, which include dust settling and the gravitational force of planets embedded in the disc. We have two main reasons to focus on M star systems. First, low-mass planets such as Super-Earths are the current preliminary targets for the ongoing and future observational missions such as Kepler. M star systems have the highest probability of detection using the transit method. In addition, some planets around M stars have already been found in or near the habitable zone. Second, Monte Carlo calculations become computationally expensive for more massive discs (such as those around classical T Tauri stars), making M star systems an ideal target. We generalize our disc models by using minimum mass solar nebula (MMSN) models in which the surface density behaves as $\Sigma\propto r^{-3/2}$, as well as models with $\Sigma \propto r^{-1}$ \citep[hereafter S07]{sjw07} used in Paper I and HP10. We then compare the type I migration of various low mass planets (2, 5, 10 $M_{\oplus }$) within our models. Thus, our paper provides a detailed exploration of the effect of the thermal structure of discs, as a result of dust settling, on planetary migration. Our plan of this paper is the following. In $\S$ \ref{disk}, we describe and compute our disc models and summarize the important points for planetary migration for two disc models with different surface density profiles. We use these disc models as a background disc structure for analytical torque calculations. In $\S$ \ref{torque}, we briefly describe the Lindblad torque and discuss the corotation torque. In $\S$ \ref{results1} and $\S$ \ref{results2}, we present our results for the cases of well mixed and dust settling, respectively. We discuss the effects of individual and combined components such as dust settling and the gravitational force of planets. The effects of the surface density profiles on the migration are also discussed. In $\S$ \ref{discus}, we discuss other heat sources and corotation torques, which are both neglected in this paper. Finally, we discuss the effects of dust settling on type II migration. In $\S$ \ref{conc}, we present our conclusions.

\label{conc} We have carefully examined Lindblad torques in radiatively heated discs. In order to make our disc models realistic, we have included dust settling and the gravitational force of planets. We have demonstrated that dust settling accelerates type I migration by up to a factor of about 2. This arises from the lower temperatures of the mid-plane region and the resultant flatter shape of discs. Both are the consequence of dust settling. We have also investigated the mass dependence on the torque, by considering planets with various masses (2, 5, 10 $M_{\oplus}$). We have found that the gravitational force of planets has two effects on the torque. One is the reduction of the torque contribution around $z=r_H$. The other is to lower the disc temperature in the mid-plane which increases the torque. In our disc models. the increment is dominant over the reduction, so that massive planets undergo even more rapid migration. We have also considered disc models with two different surface density profiles: $\Sigma \propto r^{-1}$ (S07) and $\Sigma \propto r^{-3/2}$ (MMSN). We found that, for the well mixed case, planets for the MMSN models migrate faster than those for the SO7, by up to a factor of 4. For the dust settling case, the difference becomes small (about a factor of 2). Finally, we have studied the effects of dust settling on the gap-opening mass and type II migration. The reduction of the scale height by dust settling decreases the gap-opening mass and increases the timescale of type II migration. Thus, dust settling may slightly reduce the problem of planetary migration for sufficiently massive planets. Our results show that dust settling has a clear and significant effect on planetary migration of both low and high mass planets. It is clear that an even more effective slowing mechanism is required in order to slow type I migration. Currently, two types of inhomogeneity in discs could do the job. HP10 and \citet{mpt09} demonstrated that dead zones provide two robust barriers while ice lines considered by \citet{il08v} provide another robust barrier. In our next paper, we will incorporate these barriers in one disc model, and discuss their roles in producing the mass-period relation.

1012.0254

1012

1012.0581_arXiv.txt

We report the X-ray detection of two $z>1.4$ infrared-selected galaxy clusters from the IRAC Shallow Cluster Survey (ISCS). We present new data from the {\it Hubble Space Telescope} and the W.~M.~Keck Observatory that spectroscopically confirm cluster \clB\ at $z=1.49$, the most distant of 18 confirmed $z>1$ clusters in the ISCS to date. We also present new spectroscopy for \clA, previously reported at $z = 1.41$, and measure its dynamical mass. Clusters \clB\ and \clA\ are detected in 36ks and 143ks \chandra\ exposures at significances of $5.2\sigma$ and $9.7\sigma$, from which we measure total masses of $\logmlxtwo = 14.4 \pm 0.2$ and $14.35\,^{+0.14}_{-0.11}$, respectively. The consistency of the X-ray and dynamical properties of these high redshift clusters further demonstrates that the ISCS is robustly detecting massive clusters to at least $z = 1.5$.

Present-day galaxy clusters contain large, old, roughly coeval populations of massive, quiescent galaxies spanning a vast range of local densities. As such, they provide a natural laboratory in which to test models of galaxy formation and evolution. To trace the evolution of cluster galaxies over their full lifetime, it is necessary to identify and study the {\it precursor} cluster population over a large redshift range. For instance, the Coma cluster, with a present-day mass of $\logmtwo \approx 15.3$ \citep{kubo07}, is expected to have a precursor at $z \sim 1.5$ with a halo mass of $\logmtwo \sim 14.6$. This kind of archaeology requires uniformly selected, well characterized cluster samples in which the evolutionary precursors can be statistically identified, and which are sensitive down to the group scale at very high redshift. Neither X-ray nor Sunyaev-Zel'dovich (SZ) cluster surveys have the mass sensitivity at high redshift, and optical methods fail at $z\ga 1$ as the red sequence shifts to the infrared. The \spitzer/IRAC Shallow Cluster Survey \citep[ISCS,][]{eisenhardt08} is a stellar mass-selected galaxy cluster survey spanning $0.1 < z < 2$. Clusters are identified via stellar mass overdensities in a 4.5 $\mu$m-selected galaxy sample using accurate photometric redshifts \citep{brodwin06_iss}, and their selection is therefore independent of the presence of a red sequence. There are 335 clusters and groups in the sample, identified over 7.25 deg$^2$ within the Bo\"otes field of the NOAO Deep Wide--Field Survey \citep{ndwfs99}, and 1/3 of the clusters are at $z > 1$. The ISCS cluster sample has a mean halo mass, derived from its clustering, of $\approx 10^{14}$ \msun\ out to $z=1.5$ \citep{brodwin07}. Cluster photometric redshift accuracy, based on comparison with over 100 clusters spanning $0< z< 1.5$, is excellent, with $\sigma = 0.028(1 + z)$. At $z < 1$ roughly 100 clusters have been spectroscopically confirmed, primarily using the extensive spectroscopic database of the AGN and Galaxy Evolution Survey (AGES, Kochanek \etal\ in prep). To date, we also have spectroscopically confirmed 18 clusters spanning $1 < z < 1.5$. (\citealt{stanford05,brodwin06_iss,elston06,eisenhardt08}, Brodwin \etal\ in prep; Stanford \etal\ in prep). {\it All} of the candidates for which adequate spectroscopy has been obtained have turned out to be at the predicted photometric redshifts. The ISCS is therefore the largest sample of spectroscopically confirmed galaxy clusters at $z > 1$. In this paper, we report the X-ray detection of two of the most distant ISCS clusters: \clA\ at $z=1.414$, first reported in \citet{stanford05}, and \clB\ at $z=1.487$, for which we present spectroscopic confirmation in this work. The latter is the most distant cluster to date to be spectroscopically confirmed in the ISCS. We present the spectroscopic observations and dynamical properties of these clusters in \textsection{\ref{Sec: dynamical}}. In \textsection{\ref{Sec: x-ray}} we present the X-ray observations, from which we estimate total cluster masses. In \textsection{\ref{Sec: discussion}} we compare the dynamical and X-ray properties of these clusters, including testing the well-known relation between velocity dispersion and temperature at the highest redshift yet. We also discuss the effect of possible systematic uncertainties on our measurements. We present our conclusions in \textsection{\ref{Sec: conclusions}}. We use Vega magnitudes and a WMAP7+BAO+$H_0$ $\Lambda$CDM cosmology \citep{komatsu10}: $\Omega_M = 0.272$, $\Omega_\Lambda = 0.728$, and $H_0 = 70.2$ km s$^{-1}$ Mpc$^{-1}$.

\label{Sec: conclusions} We present X-ray observations of two $z>1.4$ spectroscopically confirmed galaxy clusters from the ISCS. \clA, at $z=1.414$, is detected in deep \chandra\ data at a significance of $9.7\sigma$. We measure a luminosity of $L_{0.5 - 2} = 1.00^{+0.24}_{-0.19} \times 10^{44}~ {\rm erg}~ {\rm s}^{-1}$ and a temperature of $T_x = 3.34^{+1.90}_{-1.04}$~keV. From these we derive total $L_X$- and $T_X$-based masses of $\logmlxtwo = 14.35\,^{+0.11}_{-0.14}$ and $\logmtxtwo = 14.2\,^{+0.3}_{-0.4}$. These agree with the dynamical mass, inferred from 17 spectroscopically confirmed members, of $\logmdyntwo = 14.4\,^{+0.3}_{-0.7}$. The X-ray and dynamical properties of \clA\ are fully consistent with the $\sigma-T_X$ relation. This is the most distant cluster for which this relation has been tested. We also present spectroscopic confirmation for \clB\ at $z=1.487$, the most distant of 18 $z>1$ clusters confirmed in the ISCS to date. \clB\ is detected at a $5.2\sigma$ significance in a shallow \chandra\ exposure. We measure an X-ray luminosity of $L_{0.5 - 2} = 1.3^{+1.1}_{-0.6} \times 10^{44}~ {\rm erg}~ {\rm s}^{-1}$, from which we infer a total mass of $\logmlxtwo = 14.4 \pm 0.2$. The sparse dynamical data are consistent with this X-ray mass, confirming that \clB\ is also a massive high redshift cluster. Our primary conclusion is that the $z \ga 1.4$ clusters studied in this work, identified by the stellar-mass selection of the ISCS, are massive and have dynamical and ICM properties typical of clusters selected by the X-ray or SZE. This further confirms that stellar-mass cluster selection provides a powerful and sensitive method for studying cluster evolution to the highest redshifts. The two clusters presented in this work, with total masses of $\logmtwo \approx 14.4$ or $\approx 2-3 \times 10^{14}$ \msun, are nearly Coma--mass progenitors. The ISCS is therefore identifying the precursor population of present-day massive clusters.

1012.0581

1012

1012.5701_arXiv.txt

We analyze the production and freeze-out of the isomer \taisob\ in the $\nu$-process. We consider the influence of a possible low-lying intermediate ($J = 5$) state at 592\,keV using a transition width estimated from the measured half-life. This more realistic width is much smaller than the previous estimate. We find that the 592\,keV state leads only to a small reduction of the residual isomer population ratio from the previous result; i.e., considering this better estimate for the transition width, the isomer population ratio changes from ${\cal{R}} = 0.39$ to ${\cal{R}} = 0.38$, whereas previously it was estimated that this transition could reduce the ratio to ${\cal{R}} = 0.18$. This finding strengthens the evidence that $^{138}$La and $^{180}$Ta are coproduced by neutrino nucleosynthesis with an electron neutrino temperature of $kT \approx 4$\,MeV.

1012.5701

1012

1012.0312_arXiv.txt

Recent measurements of the Sunyaev-Zel'dovich (SZ) angular power spectrum from the South Pole Telescope (SPT) and the Atacama Cosmology Telescope (ACT) demonstrate the importance of understanding baryon physics when using the SZ power spectrum to constrain cosmology. This is challenging since roughly half of the SZ power at $\ell$=3000 is from low-mass systems with $10^{13} h^{-1}$ M$_{\odot} < M_{500} < 1.5\times10^{14} h^{-1}$ M$_{\odot}$, which are more difficult to study than systems of higher mass. We present a study of the thermal pressure content for a sample of local galaxy groups from Sun et al.~(2009). The group $Y_{\rm sph, 500} - M_{500}$ relation agrees with the one for clusters derived by Arnaud et al.~(2010). The group median pressure profile also agrees with the universal pressure profile for clusters derived by Arnaud et al.~(2010). With this in mind, we briefly discuss several ways to alleviate the tension between the measured low SZ power and the predictions from SZ templates.

Galaxy groups are not scaled-down versions of rich clusters following simple self-similar relations (e.g., Ponman et al. 2003; Voit 2005). They are systems where the role of complex baryon physics (e.g., cooling, galactic winds, and AGN feedback) begins to dominate over gravity. The effects of these baryon processes are not large but still significant in massive clusters, and therefore need to be calibrated if we want to further improve the cosmological constraints from clusters. Since the role of these processes is less pronounced in massive clusters, it is easier to study and understand them by observing groups. The importance of galaxy groups for cosmology has also been well demonstrated by the recent measurements of the Sunyaev-Zel'dovich (SZ) angular power spectrum. SPT measured a value for the SZ power at $\ell=3000$ (scales of $\sim 4'$) that was lower than most prior predictions by at least a factor of two (Lueker et al. 2010). A measurement of the SZ power spectrum by ACT is consistent with these results (Das et al.~2010; Dunkley et al. 2010). The SZ angular power spectrum is a sensitive probe both of cosmological parameters and of the hot gas content of galaxy clusters and groups (e.g. Komatsu \& Seljak 2002). Regarding the latter, it opens a new window into low-mass, high-redshift systems as about half of the SZ power at $\ell$=3000 comes from halos with $10^{13} h^{-1}$ M$_{\odot} < M_{500} < 1.5\times10^{14} h^{-1}$ M$_{\odot}$ and $z > 0.5$ (e.g., Trac et al.~2010). While the examination of the thermal pressure content in $z > 0.5$ groups is a challenge to current X-ray telescopes with typical exposures, such work can be done for local groups. Group pressure profiles have received little attention to date and the existing samples are small (e.g., Mahdavi et al. 2005; Finoguenov et al. 2007). In this letter, we present the pressure profiles of hot gas in 43 local galaxy groups from the Sun et al. (2009, S09 hereafter) sample. We assume $\Omega$$_{\rm M}$=0.24, $\Omega_{\rm \Lambda}$=0.76 and H$_{0}$ = 73 km s$^{-1}$ Mpc$^{-1}$.

The results above suggest that the thermal pressure of local galaxy groups from the S09 sample is consistent with the extrapolation from the A10 results, although statistical errors and scatter are still large. Interestingly, recent measurements of the SZ angular power spectrum are at least a factor of two lower than prior expectations at $\ell=3000$ (Lueker et al.~2010). The thermal SZ power spectrum scales roughly as the square of the thermal SZ flux. Given the results presented above, we briefly discuss several possibilities that may alleviate this tension. {\bf X-ray selection bias:} The S09 sample is an X-ray archival sample and the REXCESS sample is an X-ray-luminosity-selected sample. Both samples can be different from mass-selected samples. The \chandra\ archive may be biased to systems with bright cores, while X-ray under-luminous groups and clusters may exist (e.g., Rasmussen et al. 2006; Popesso et al. 2007). However, the $Y_{\rm sph} - M$ relation at $r_{500}$ and beyond is less affected by the presence of X-ray bright regions (e.g., a large cool core) than the $L_{\rm X} - M$ relation. For the S09 sample, 12\% - 68\% of the X-ray flux (a median of $\sim$ 34\%) is from within $0.15 r_{500}$, while such regions only contribute $\sim$ 5\% to $Y_{\rm sph, 500}$ for $M_{500} = 10^{13} h^{-1}$ - $10^{15} h^{-1}$ M$_{\odot}$ halos, assuming the A10 pressure profile. The contribution may be even smaller than 5\% for X-ray under-luminous systems as their gas cores are fainter than those of the REXCESS clusters used to derive the A10 profile. One main conclusion of the S09 work is that the gas content of groups is comparable to that of clusters at $r > r_{2500}$, at least for the S09 sample. If we combine the $n_{\rm e} - T_{500}$ relations from Vikhlinin et al. (2009), S09 and REXCESS, the trend of slope flattening with increasing radius is significant, with an almost constant density at $r_{500}$ from groups to clusters. This trend is consistent with the scenario that much of the low-entropy gas in low-mass systems has been ejected to large radii by strong feedback (e.g., McCarthy et al. 2010). However, it remains to see whether this result applies to mass-selected samples. One way to test this is to examine scaling relations from non-X-ray-selected samples. This kind of work has been done on optically selected samples by stacking the \rosat\ all-sky survey data (Dai et al. 2010; Rykoff et al. 2008). Besides the systematic uncertainties with stacking, the \rosat\ temperatures from stacking are often biased (Rykoff et al. 2008) and contamination to those samples, especially at the low-mass end, can be severe. {\bf Pressure contribution at ${\bf r > r_{500}}$:} The total SZ flux is more sensitive to the gas in cluster outskirts ($r > r_{500}$) than the total X-ray flux. Few direct X-ray constraints exist at such large radii, especially for groups. Although the contribution from $r > r_{500}$ to $Y_{\rm sph}$ is significant (e.g., 40\% - 70\% increase by integrating to 2$r_{500}$ for the A10 profile), the contribution to the SZ power spectrum at $\ell$=3000 assuming the A10 profile is smaller, only about 20\% from $r > r_{500}$ regions. So an overestimate of the thermal pressure from the A10 model only at $r > r_{500}$ could overpredict the SZ power spectrum by at most 20\% at $\ell$=3000. {\bf Dynamical state of the ICM:} The SZ signal measures the total thermal energy of electrons. However, the potential energy of halos may not be fully converted into thermal energy of electrons because of e.g., recent mergers and weak viscosity of the ICM (e.g., Lau et al. 2009; Burns et al. 2010). Galaxies can also contribute to the non-thermal pressure support by e.g., injected magnetic fields and cosmic rays. The non-thermal pressure support may also cause the ICM to be clumpy. For a clumpy ICM, the SZ signal predicted from the X-ray data will be biased high. All these effects may have a dependence on mass (or the ICM temperature), and the evolution of these effects with redshift may not be self-similar. The impact of the non-thermal pressure support on the SZ power spectrum has been discussed by Shaw et al. (2010), who examined models with a radial dependence of the non-thermal pressure. Trac et al. (2010) examined a model with $20\%$ non-thermal pressure for all clusters and groups at all masses and redshifts. However, both models do not predict the mean value measured for the SZ power spectrum at $\ell=3000$ by Lueker et al. (2010), being high by about 1$\sigma$. Interestingly, SZ observations suggest the latter model predicts too little SZ flux for very massive clusters (Sehgal et al.~2010b). As for clumpiness, the good agreement between the measured SZ radial profile and the prediction from X-ray data for individual clusters (e.g., Plagge et al 2010; Komatsu et al. 2010; Sehgal et al. 2010b) suggests that clumpiness should be weak for massive clusters. However, one can imagine a mass dependence for clumpiness, as both heat conduction and dynamic viscosity can be much weaker in groups than in clusters. If the hydrostatic equilibrium mass under-estimates the true mass by 20\% from $M_{500}$ = 10$^{13}$ h$^{-1}$ M$_{\odot}$ - 10$^{15}$ h$^{-1}$ M$_{\odot}$ (e.g., Nagai et al. 2007), $Y_{\rm sph, 500}$ will increase by 4\% - 7\%, if the universal pressure profile from A10 is assumed. Therefore, the normalization of the $Y_{\rm sph, 500} - M_{500}$ relation will decrease by $\sim$ 24\%, which roughly translates to $\sim$ 42\% decrease on the predicted SZ power spectrum. If the mass bias is larger for low-mass halos, the relation will be steeper and the decrease will be larger. However, it is a big challenge to constrain the effects of non-thermal pressure support in groups and clusters, especially its dependence with mass and redshift. For groups, robust mass measurements that do not assume hydrostatic equilibrium are required. Two promising methods are stacking of the lensing data (e.g., Leauthaud et al. 2010) and caustics (Rines \& Diaferio 2010). Future X-ray microcalorimeter observations (e.g., by {\em Astro-H}) may also constrain the ICM turbulence directly. {\bf Evolution of the ICM properties:} While local groups are discussed in this paper, most of the SZ power at $\ell$=3000 from groups is from systems at $z>0.5$. Evolution of the ICM properties is poorly constrained for poor clusters and groups. Recent results on $z \geq 0.5$ groups suggest that the evolution of the $L_{\rm X}$ scalings is not weaker than the self-similar prediction (e.g., Jeltema et al. 2009; Leauthaud et al. 2010), but the statistical uncertainties are large. Proper understanding of the evolution of the $L_{\rm X}$ scalings requires a good understanding of the selection function (e.g. Pacaud et al. 2007). Better constraints on the evolution of the low-mass end of the $L_{\rm X}$ scaling relations should be achieved with more \xmm\ and \chandra\ data on larger samples with well-defined selection functions. Of course, the evolution of the $L_{\rm X}$ scaling relations is not equal to the evolution of the $Y - M$ relation. More factors, e.g., the evolution of the cool core fraction and the gas distribution, need to be accounted for. Alternatively, deep SZ observations of high-$z$ groups, either individually or by stacking, can directly constrain the evolution of the $Y - M$ relation. {\bf Contamination from radio and infrared galaxies:} Both radio and infrared galaxies could potentially fill in SZ decrements at 150 GHz. While the contamination from radio galaxies should be small (see discussion in Sehgal et al 2010a), the contamination from dusty star-forming (infrared) galaxies is less clear. In Lueker et al. (2010), the signal from infrared galaxies was removed from maps at 150 GHz by subtracting maps at 220 GHz after fitting for a weighting factor. If all infrared sources have the same spectral index of $\alpha = 3.6$, then this should effectively remove infrared contamination from the 150 GHz maps. If some infrared sources have a shallower slope (e.g., $\alpha = 2.6$ as in Knox et al. 2004), then residual contamination will remain. However, a more recent analysis by Shirokoff et al. (2010) suggests that even a large correlation between infrared galaxies and groups/clusters would not increase the 95\% CL upper limit on the thermal SZ power spectrum to the level that it is consistent with predictions prior to Lueker et al (2010). \\ This work shows that the local groups from the S09 sample follow the extrapolation of the pressure scaling relations from A10. More data are required to reduce the statistical errors and more importantly explore the systematic uncertainties discussed above. Regarding the low SZ power measured by recent experiments, we suggest some astrophysical possibilities that may alleviate the apparent tension between models and measurements. Understanding the SZ power spectrum will provide important insights into both baryon physics and cosmology.

1012.0312

1012

1012.2121_arXiv.txt

We report the discovery of a \period\ X-ray pulsar in observations of the soft $\gamma$-ray source \igr\ with the {\it Rossi X-ray Timing Explorer} (\xte). \psr\ is spinning down rapidly with period derivative \pdotl, yielding a spin-down luminosity \edot, characteristic age \tauc, and surface dipole magnetic field strength \bs. Within the \integral\ error circle lies a point-like \xmm\ and \chandra\ X-ray source that shows evidence of faint extended emission consistent with a pulsar wind nebula (PWN). The \xmm\ spectrum of the point source is well fitted by an absorbed power-law model with photon index $\Gamma_{PSR} = 1.1 \pm 0.2$, $N_{\rm H} = (4.3 \pm 0.6) \times 10^{22}$~cm$^{-2}$, and $F_{\rm PSR}(2-10\,{\rm keV}) = (3.8 \pm 0.3) \times 10^{-12}$ erg~cm$^{-2}$~s$^{-1}$, while the spectral parameters of the extended emission are roughly $\Gamma_{\rm PWN} \approx 2.1$ and $F_{\rm PWN}(2-10\,{\rm keV}) \approx 9 \times 10^{-13}$ erg~cm$^{-2}$~s$^{-1}$. \igr\ is also coincident with the compact TeV source \tev. For an assumed distance of 7~kpc in the Scutum arm tangent region, the $0.35-10$~TeV luminosity of \tev\ is 0.13\% of the pulsar's spin-down energy, while the ratio $F(0.35-10\,{\rm TeV})/F_{\rm PWN}(2-10\,{\rm keV}) \approx 2$. These properties are consistent with leptonic models of TeV emission from PWNe, with \psr\ in a stage of transition from a synchrotron X-ray source to an inverse Compton $\gamma$-ray source.

The detection of $10^{12}$~eV emission associated with supernova products in the Galaxy has opened up a new window on the evolution of these energetic stellar remnants. More than 2/3 of the now $>60$ Galactic TeV sources are supernova remnants or pulsar wind nebulae (PWNe), the latter being the largest class\footnote{VHE $\gamma$-ray Sky Map and Source Catalog,\\ \url{http://www.mppmu.mpg.de/$\sim$rwagner/sources/index.html}}. Many of these TeV PWNe are spatially offset from middle-aged ($\tau_c \sim 10^{4}-10^{5}$~year old) pulsars, and are often more extended and more luminous than their X-ray PWNe counterparts. This TeV emission is likely to be inverse Compton scattered radiation from relic electrons produced by the pulsar in an earlier stage of energetic spin-down \citep{dej08,zha08,mat09b} In contrast, younger ($\tau_c \sim 10^3$~year old) pulsars are associated with compact TeV sources that are co-located with their X-ray PWNe. In these younger systems the high magnetic fields make for efficient synchrotron X-ray sources, and inefficient inverse Compton TeV emission. \tev\ is one of the fainter sources detected in the HESS Galactic Plane survey \citep{aha05,aha06} and subsequent dedicated observations, with significance of $6.4\sigma$, flux $>350$~GeV of $\approx 15$~mCrab, and only weak indication of extent \citep{ter08}. \tev\ is coincident in position with the \integral\ source \igr\ that was discovered during a survey of the tangent regions of the Sagittarius and Scutum spiral arms \citep{mol04,bir06}. Follow-up X-ray observation of \igr\ by \citet{rod08} with the \swift\ X-ray telescope (XRT) located a highly absorbed X-ray point source within the \integral\ error circle. These authors suggested a Galactic X-ray binary origin for \igr\ based on a $K$-band Two Micro All Sky Survey star located within the \swift\ XRT error circle. However, that association was excluded by the precise X-ray position obtained in a brief \chandra\ High Resolution Camera (HRC) observation \citep{rat10}. \citet{ter08} used imaging and spectroscopic evidence from an \xmm\ observation to argue that \igr\ is instead a pulsar/PWN. As such, it would join the dozen hard X-ray members of this class detected by \integral\ \citep{mat09a,ren10}. In Section~2 we present in detail the \xmm\ imaging and spectral data that support a pulsar/PWN interpretation for \igr. In Section~3, we report the discovery using {\it Rossi X-ray Timming Explorer} (\xte) of \psr\ and its spin-down, which verifies this conjecture. In Section~4, we discuss the properties of \tev\ in the context of the spin-down parameters of \psr.

The discovery of \psr\ using \xte\ verifies the conjecture of \citet{ter08} that \xmmu, \igr, and \tev\ are all manifestations of a young pulsar/PWN system. Even though the \xte\ field of view is wider than those of \xmm\ and \chandra, there is little doubt that \psr\ is the compact source in \igr, given the morphological and spectral evidence, and the compatibility of the \xte\ measured pulsed flux with that of \xmmu. \citet{ter08} estimated that the spin-down luminosity of \psr\ would be $\approx 9 \times 10^{36}$~erg~s$^{-1}$ based on the empirical correlation between $\dot E$ and $\Gamma_{\rm PSR}$ of \citet{got03}; this turns out to have been an accurate prediction. Assuming a distance of 7~kpc as proposed in Section 2.1, the X-ray luminosities of the pulsar and PWN, $L_{\rm PSR}(2-10\,{\rm keV}) = 2.9 \times 10^{-3}\,\dot E\,d_{7}^2$ and $L_{\rm PWN}(2-10\,{\rm keV}) = 6.6 \times 10^{-4}\,\dot E\,d_{7}^2$, respectively, are consistent with the range of pulsar/PWN systems \citep{kar08}. The $20-100$~keV flux of \igr\ measured by \integral\ \citep[$2 \times 10^{-11}$ erg~cm$^{-2}$~s$^{-1}$;][]{ter08} corresponds to $L(20-100\,{\rm keV}) = 0.012\,\dot E\,d_{7}^2$. Given the $2-20$~keV flux and spectral information from \xmm\ and \xte, the \integral\ source is likely to be dominated by the pulsar rather than the PWN. The TeV flux from \tev\ \citep[$2.2 \times 10^{-12}$ erg~cm$^{-2}$~s$^{-1}$;][]{ter08} corresponds to $L(0.35-10\,{\rm TeV}) = 1.3 \times 10^{-3}\,\dot E\,d_{7}^2$. Such a small efficiency of converting spin-down power to high-energy radiation is typical of high $\dot E$ pulsars. The ratio $F(0.35-10\,{\rm TeV})/F_{\rm PWN}(2-10\,{\rm keV}) \approx 2$ falls squarely on the inverse correlation between this quantity and $\dot E$ that was fitted by \citet{mat09b}, and modeled by them in terms of an evolving PWN emitting synchrotron X-rays and inverse Compton $\gamma$-rays. \psr\ is in transition from a synchrotron dominated X-ray PWN to an inverse Compton scattered TeV nebula, an expected phase through which a middle-aged pulsar will pass.

1012.2121

1012

1012.3194_arXiv.txt

{Millisecond pulsars are believed to be old pulsars spun up by a surrounding accretion disc. Magnetic fields are thought to play a leading role in this, both by determining the location of the inner edge of the disc and by exerting an additional torque on the star (as a result of the interaction between the stellar magnetic field and the disc plasma motion, which creates a toroidal component $B_\phi$). In some well-known analytic models, developed in the 1980s, the $B_\phi$ profile was taken to be proportional to the relative angular velocity between the disc plasma and the neutron star, multiplied by a vertical dipolar field. The present work stands in the line of improving those models, suggesting a new profile for $\mathbf{B}$. In a previous paper, we discussed the poloidal component of the magnetic field and here we consider the toroidal component, again making the kinematic approximation and looking for steady solutions of the induction equation for axisymmetric models. The poloidal magnetic field is not assumed to be dipolar and the poloidal velocity field is not taken to be zero everywhere. We also do not use the thin disc approximation to simplify the induction equation but instead solve it numerically in full 2D. The profile obtained in the earlier analytic models is shown to have very limited validity and a more general semi-analytic solution is proposed.}

\label{sec:INTRO} In the present work we study the deformation caused in a neutron star's magnetic field because of the interaction with the matter in a surrounding accretion disc. A basic description for this kind of system was given by Ghosh and Lamb in \cite{GL79a}, with the model subsequently being improved by Wang (\cite{W87}) and Campbell (\cite{C87}), who suggested an analytic expression for the toroidal component of the field. This expression has been widely used since then, on account of it being both simple and physically plausible. In these analytic models, the authors made the kinematic approximation, looking for an axisymmetric stationary solution of the induction equation with a given unchanging structure for the fluid in the disc. They further took the disc to be thin, the poloidal component of the magnetic field to be exactly dipolar, and the velocity field to have zero poloidal component, with its azimuthal component being Keplerian inside the disc\footnote{Campbell also considered non-Keplerian flow in the inner part of the disc.} and corotating at the top of the disc. Using cylindrical coordinates {($\varpi$, $\phi$, $z$)}, they found \begin{equation} \label{eq:bp_an} B_\phi = \gamma_{\rm a} \, (\Omega_K - \Omega_{\rm s}) \, B_z \, \tau_{\rm d} \propto \Delta \Omega / \varpi^3 \, {\rm ,} \end{equation} where $\gamma_{\rm a}$ is the amplification factor, $\Omega_K$ and $\Omega_{\rm s}$ are the Keplerian and stellar angular velocities, respectively, and $\tau_{\rm d}$ is the dissipation time scale. The amplification factor $\gamma_{\rm a}$ was taken to be a constant not much greater than $1$ (it depends on the steepness of the transition -- in the z-direction -- between Keplerian motion inside the disc and corotation with the star at the top of the disc). The precise profile of $\tau_{\rm d}$ depends on what is the dominant mechanism for dissipating the magnetic field. Wang (\cite{W95}) considered three different cases, with $\tau_{\rm d}$ being dominated by the Alfven velocity, turbulent diffusion and magnetic reconnection, respectively. Equation (\ref{eq:bp_an}) was then used for calculating the net magnetic torque exerted on the neutron star. The vertically averaged torque can be written as \begin{equation} \label{eq:tor} \bar{\Gamma} \propto B_\phi \, B_z \, \varpi / h \, {\rm ,} \end{equation} where $h$ is the semi-thickness of the disc. Regions of the disc inward of the corotation point therefore give positive contributions to the torque (because $B_\phi>0$), while the remainder of the disc gives negative contributions (because $B_\phi<0$). The total magnetic torque is obtained by integrating the local values from the inner edge to the outer boundary, and it can be either positive or negative depending on the location of the inner edge of the disc with respect to the corotation point. The aim of the present paper is to develop a semi-analytic model that can improve on those of Wang and Campbell, while remaining simple enough to be useful for people discussing the behaviour of astrophysical sources, giving a conceptual picture to go alongside results from large numerical calculations where the full set of the MHD equations is solved. \begin{figure}[ht] \centering \includegraphics[width=0.4\textwidth]{figures/16314fg1.eps} \caption{Schematic representation of our model (not to scale). We use $r_{\rm in} = 10\,r_{{\rm g}}$, $r_{\rm tr} \sim 22 \, r_{{\rm g}}$ and $r_{\rm lc} \sim 115\,r_{{\rm g}}$. The opening angles are $8^{\circ}$ for the disc alone and $10^{\circ}$ for the disc plus corona. The outer disc extends much further out than the main region shown here: the grid continues until $r_{\rm out} = 380 \, r_{{\rm g}}$.} \label{fig:model} \end{figure} Our approach, for the time being, is to continue to retain axisymmetry and the kinematic approximation but to calculate a consistent steady-state solution for the magnetic field, relaxing the assumptions on the poloidal components of the magnetic and velocity fields and using a 2D model without any vertical averaging of the Taylor expansion of the induction equation. In the main region of the outer disc (see Fig. \ref{fig:model}) we use a simple Keplerian velocity profile, but this is something that will be improved on later. In a previous paper (Naso \& Miller \cite{papI}, hereafter Paper I) we analysed the distortion of the poloidal component of the magnetic field using a similar approach, and found that deviations away from a dipole field can be quite significant. Here we focus on the toroidal component and use the results of the previous model to solve the $\phi$ component of the induction equation. We find that in general $B_\phi$ follows a profile different from that of the analytic models, i.e. Eq. (\ref{eq:bp_an}), and reduces to that only in a very particular case. Following this introduction, in Sect. \ref{sec:MOD} we briefly describe our model, which is the same as that of Paper I; in Sect. \ref{sec:EQU} we recall the equations used (obtained from the induction equation), give expressions for the velocity and diffusivity profiles and outline our solution method (details of tests made on the code are given in an Appendix); in Sect. \ref{sec:RES} we present our numerical results; in Sect. \ref{sec:DIS} we comment on these, comparing them with those of the earlier analytic models, and develop our new suggestion for the $B_\phi$ profile. Sect. \ref{sec:CON} contains conclusions.

\label{sec:CON} In this paper we have considered a system consisting of a rotating neutron star, having a dipole magnetic field aligned with the rotation axis and surrounded by an accretion disc. The disc is truncated at the Alfven radius and has a coronal layer above and below it. The region outside the corona is taken to be vacuum and we impose dipolar boundary conditions on all of the boundaries. Our aim was to improve on the analytic models developed by Wang (\cite{W87}) and Campbell (\cite{C87}) (W\&C). As in those models, we have made the kinematic approximation and have looked for an axisymmetric stationary solution of the induction equation, but we have gone beyond those models in solving for all of the components of the magnetic field and not assuming the poloidal component to be dipolar within the disc. We have also retained all of the components of the velocity field rather than putting $v_r$ and $v_\theta$ to zero everywhere. We have performed a fully two-dimensional calculation, without making any vertical average or Taylor expansion in $h$ (the semi-height of the disc). Finally we have neglected dynamo action but have included a turbulent magnetic diffusivity. The analysis of the poloidal component of the magnetic field has been presented in a previous paper (Naso \& Miller, \cite{papI}, Paper I); in the present paper we have focused on the toroidal component. We have solved the $\phi$ component of the induction equation numerically and have shown that the profile obtained for $B_\phi$ can be very different from that in the earlier analytic models. In the W\&C models, the toroidal field strength was taken to be proportional to the relative angular velocity between the disc and the central object multiplied by the vertical field, which was taken to be dipolar. However in Paper I we found that, when calculated consistently, the poloidal field component was often far from being dipolar. Therefore a first improvement with respect to the earlier models was to use the poloidal field as obtained in our calculations, i.e. a field dragged inwards by the plasma motion. This behaviour explains why we then find different intensities for the toroidal field, and also a different location for its global minimum. When the turbulent magnetic diffusivity $\eta$ increases or the radial velocity decreases one expects the field to be progressively less distorted by the plasma motion. This is indeed what we have found both here and in Paper I. Our results show that when the diffusivity $\eta_0$ is larger than about $10^{12}$ cm$^2$ s$^{-1}$, with the characteristic velocity $v_0$ being of the order of $10^5$ cm s$^{-1}$, then the field is barely modified. Therefore whenever we expect the magnetic field to deviate from the stellar dipole, $\eta$ should not be larger than $10^{7} \, |v_{\rm cgs}|$ cm$^2$ s$^{-1}$, where $|v_{\rm cgs}|$ is the characteristic magnitude of the radial velocity expressed in cm s$^{-1}$. When the turbulent magnetic diffusivity is not constant throughout the disc and corona, two additional extremal points may well appear: if the radial derivative of the magnetic distortion function $D_{\rm m}$ is larger than a critical value (about $0.04$ in the equatorial plane), there is an additional maximum and minimum, and in some cases $B_\phi$ can even become positive again outward of the corotation point, so that there are additional locations where $B_\phi=0$. It is clear that under these circumstances the picture of which regions of the disc tend to spin the star up or down has to be radically redrawn (this will be the subject of a future investigation). However we should emphasise here that there are still many uncertainties among experts about which profile of $\eta$ should be used and we have therefore made very simple choices here in line with our step-by-step approach. We have presented a new suggestion for the $B_\phi$ profile, which reduces to that of W\&C if one imposes $B_r = \partial_\theta B_\theta = 0$ and $\Omega = \Omega_K$. In general there are large parts of the disc where the additional terms included in our new picture for $B_\phi$ dominate over the one retained by W\&C (see Fig.~\ref{fig:RG}). Our simplified expression (Eq.~\ref{eq:my_bp}) reproduces the numerical results quite well (compare Figs.~\ref{fig:ref_cont_10.115} and \ref{fig:a0_10.115}), the differences being due to approximations made in calculating the generation and loss terms for $B_\phi$. Summarising, in our calculations we have found that $B_\phi$ can have two maxima and two minima (see Fig.~\ref{fig:ref_cont_10.115}). The first maximum (positive and inward of the corotation point) and the first minimum (negative and outward of the corotation point) can be explained referring to the quantity $\Delta \Omega \, B_\theta$, which has two extrema at the same locations as for $B_\phi$ (see Fig.~\ref{fig:DWBt_ref}). These extrema appear also in the W\&C models, where the toroidal field is, in fact, taken to be proportional to $\Delta \Omega \, B_z$. There is a fundamental difference however: in W\&C $B_z$ is a pure dipole, whereas the $B_\theta$ which we consider here is that of a field being dragged inward by the motion of the accreting material. When $\eta$ is not constant, there is an additional maximum whose magnitude and sign depend on the diffusivity in the disc, and whose radial location is always outward of the first minimum, coinciding with that of the maximum in the radial derivative of the magnetic distortion function $D_{\rm m}$. Outward of this maximum, the field tends to come back to the profile that it would have had if $\eta$ were constant, and this produces the last minimum (compare Figs.~\ref{fig:ref_eta_avgs} and \ref{fig:eta_const_av}). The main conclusion of this analysis is that, when the poloidal component of the magnetic field is treated self-consistently in the calculations, the profile for $B_\phi$ can be significantly different from that obtained by W\&C, and the magnetic torque generated by it would then be different as well. Moreover, when the turbulent magnetic diffusivity $\eta$ is not constant throughout the disc and corona, some additional unexpected features can appear (such as a region of positive $B_\phi$ outward of the corotation point). In the present work, we have retained the very simple Keplerian rotation law in the main part of the disc. Even within a purely hydrodynamical treatment, more complicated velocity fields than this are expected (see Kluzniak and Kita, \cite{KK2000}; Jiao and Wu, \cite{JW11}) and further changes are expected when back-reaction from the magnetic field on the velocity field is included. The effects of this will be another topic for investigation in subsequent stages of our step-by-step approach. \appendix

1012.3194

1012

1012.4582_arXiv.txt

The ultraviolet stellar wind lines of the photometrically periodic variable early B-type star $\sigma$ Lupi were found to behave very similarly to what has been observed in known magnetic B stars, although no periodicity could be determined. AAT spectropolarimetric measurements with SEMPOL were obtained. We detected a longitudinal magnetic field with varying strength and amplitude of about 100 G with error bars of typically 20 G. This type of variability supports an oblique magnetic rotator model. We fold the equivalent width of the 4 usable UV spectra in phase with the well-known photometric period of 3.019 days, which we identify with the rotation period of the star. The magnetic field variations are consistent with this period. Additional observations with ESPaDOnS attached to the CFHT strongly confirmed this discovery, and allowed to determine a precise magnetic period. Like in the other magnetic B stars the wind emission likely originates in the magnetic equatorial plane, with maximum emission occurring when a magnetic pole points towards the Earth. The 3.0182 d magnetic rotation period is consistent with the photometric period, with maximum light corresponding to maximum magnetic field. No helium or other chemical peculiarity is known for this object.

1012.4582

1012

1012.2647_arXiv.txt

Inflation (for the general review, \cite{textbook}) solves many fine tuning problems of conventional hot big bang model and provides the seeds of structure formation. The spectrum of the primordial curvature perturbation is given by \begin{equation} P_R=\frac{1}{12\pi^2M_P^6}\frac{V^3}{V'^2} \;. \label{spectrum} \end{equation} The number of e-folds is given by \begin{equation} N=M^{-2}_P\int^{\phi(N)}_{\phi_{end}}\frac{V}{V'}d\phi. \label{efolds} \end{equation} The spectrum corresponds to the CMB scale which we can observe left horizon at $N \sim 60$ and the spectrum is normalized to be $P_R^{1/2} = 5 \times 10^{-5}$ by the Cosmic Microwave Background (CMB) temperature fluctuations. The latest WMAP result \cite{Komatsu:2010fb} suggests a red spectrum with the spectral index $n_s \sim 0.96$. Future experiments like PLANCK satellite may further confirm this result. This talk is organized as follows. In section~\ref{1}, I introduce the idea of hybrid inflation. In section~\ref{2}, a particular realization of hybrid inflation based on supersymmetry called supernatural inflation is introduced. I also explain the current problem of the model compared with experiments. In section~\ref{3}, I explain the idea of hilltop supernatural inflation and how to realize it. In section~\ref{4}, I will show the constraint to the reheating temperature as a function of gravitino mass. Section~\ref{5} is my conclusion.

\label{5} In this talk I have already shown that it is possible to reduce the spectral index to the value suggested by current observation. We assume SUSY breaking is gravity mediated. The cosmological consequences of gauge mediation has been considered in another work \cite{Lin:2009ux}.

1012.2647

1012

1012.0642_arXiv.txt

{ Using the 32-m Medicina, 45-m Nobeyama, and 100-m Effelsberg telescopes we found a statistically significant velocity offset $\Delta V \approx 27 \pm 3$ \ms\ $(1\sigma)$ between the inversion transition in \nhhh (1,1) and low-$J$ rotational transitions in \nnhp (1-0) and \hcccn (2-1) arising in cold and dense molecular cores in the Milky Way. Systematic shifts of the line centers caused by turbulent motions and velocity gradients, possible non-thermal hyperfine structure populations, pressure and optical depth effects are shown to be lower than or about 1 \ms\ and thus can be neglected in the total error budget. The reproducibility of $\Delta V$ at the same facility (Effelsberg telescope) on a year-to-year basis is found to be very good. Since the frequencies of the inversion and rotational transitions have different sensitivities to variations in $\mu \equiv m_{\rm e}/m_{\rm p}$, the revealed non-zero $\Delta V$ may imply that $\mu$ changes when measured at high (terrestrial) and low (interstellar) matter densities as predicted by chameleon-like scalar field models~-- candidates to the dark energy carrier. Thus we are testing whether scalar field models have chameleon-type interactions with ordinary matter. The measured velocity offset corresponds to the ratio $\Delta \mu/\mu \equiv (\mu_{\rm space} - \mu_{\rm lab})/\mu_{\rm lab}$ of $(26 \pm 3)\times10^{-9}$ ($1\sigma$). }

\label{sec:1} This contribution sums up our results of differential measurements of the electron-to-proton mass ratio, $\mu = m_{\rm e}/m_{\rm p}$, carried out at the 32-m Medicina, 45-m Nobeyama, and 100-m Effelsberg telescopes \cite{LMK08, MLK09, L10a, L10b}. With high spectral resolution (FWHM $\sim$30-40 \ms), we observed narrow emission lines (FWHM $< 200$ \ms) of N-bearing molecules arising in cold and dense molecular cores of starless molecular clouds located in the Milky Way. The cores are characterized by low kinetic temperatures $T_{\rm kin} \sim 10$ K, gas densities $n \sim 10^4 - 10^5$ \cmm, magnetic fields $B < 10$ $\mu$G, and ionization degrees $x_e \sim 10^{-9}$ \cite{Di07}. The objective of this study is to probe the value of $\Delta \mu/\mu \equiv (\mu_{\rm space} - \mu_{\rm lab})/\mu_{\rm lab}$ which is predicted to be variable when measured at high (laboratory) and low (space) matter density environments \cite{OP08}. A hypothetical variability of $\mu$ is thought to be due to the scalar fields~-- candidates to the dark energy carrier,~-- which are ultra-light in cosmic vacuum but possess an effectively large mass locally when they are coupled to ordinary matter by the so-called chameleon mechanism \cite{KW04}. Several possibilities to detect chameleons were discussed in \cite{Bur09,Dav09}. First laboratory experiments constraining these models have been recently carried out at Fermilab \cite{U10} and in the Lawrence Livermore National Laboratory \cite{Ry10}. A subclass of chameleon models considers the couplings of a scalar field to matter much stronger than gravitational and predicts that fundamental physical quantities such as elementary particle masses may depend on the local matter density, $\rho$, \cite{OP08}. This means that a nonzero value of $\Delta \mu/\mu$ is to be expected for all interstellar clouds irrespective of their position and local matter density because the difference $\Delta \rho$ between the terrestrial environment in laboratory measurements and dense interstellar molecular clouds is extremely large, $\rho_\oplus/\rho_{\rm cloud} > 10^{10}$. In the standard model (SM) of particle physics the dimensionless mass ratio $\mu = m_{\rm e}/m_{\rm p}$ defines the ratio of the electroweak scale to the strong scale since the mass of the electron $m_{\rm e}$ is proportional to the Higgs vacuum expectation value and the mass of the proton $m_{\rm p}$ is proportional to the quantum chromodynamics scale $\Lambda_{\rm QCD}$ \cite{CFK09}. The SM is extremely successful in explaining laboratory physics, but it has serious problems in astrophysics where it completely fails to explain dark matter and dark energy. There are many extensions of the SM including supersymmetry and different multidimensional theories which introduce new particles as possible candidates for the dark matter and additional scalar fields to describe the nature of dark energy. A concept of dark energy with negative pressure appeared in physics long before the discovery of the accelerating universe through observations of nearby and distant Type Ia Supernovae \cite{Per98, Ri98}. Early examples of dark energy in the form of a scalar field with a self-interaction potential can be found in reviews \cite{Ca98} and \cite{PR03}. If masses of the elementary particles are affected by scalar fields, one can probe the dimensionless constant $\mu$ at different physical conditions by means of high precision spectral observations. Namely, since the inversion and rotational molecular transitions have different sensitivities to variations in $\mu$ \cite{FK07}, a nonzero $\Delta \mu/\mu$ causes an offset between the radial velocities ($\Delta V \equiv V_{\rm rot} - V_{\rm inv}$) of \textit{co-spatially} distributed molecules, which, in turn, provide a measure of $\Delta \mu/\mu$. When the inversion line belongs to \nhhh we will herein call this procedure the ammonia method.

\label{sec:6} In two molecular cores with the lowest Doppler noise L1498 and L1512, we register very close values of the velocity offset \DV\ $\sim 27$ \ms\ between the rotational transition \hcccn(2--1) and the inversion transition \nhhh(1,1). These values coincide with the most accurate estimate obtained from the Effelsberg dataset on 12 molecular clouds in the Milky Way \cite{L10a}. Two other cores, L1517B and L1400K, exhibit velocity shifts that are either higher ($\sim47$ \ms\ in L1517B) or lower ($\sim 9$ \ms\ in L1400K) than the mean value, but the positive (L1517B) and negative (L1400K) deflections from the mean can be explained from the observed kinematics in these cores. Of course, simultaneous observations of the \nhhh(1,1) inverse transition and rotational transitions of some N-bearing molecules, such as \nnhp(1-0) and \nndp(1-0) would give a more accurate test. The main obstacle to this way is that the laboratory frequencies of \nnhp\ and \nndp\ are known with accuracies not better than 14 \ms\ \cite{L10a}. Using the \nnhp\ rest frequency from the Cologne Database for Molecular Spectroscopy (CDMS) \cite{Mu05} and observing with the Nobeyama 45-m telescope, we obtained a velocity shift between \nnhp\ and \nhhh\ of $23.0\pm3.4$ \ms\ in L1498, $24.5\pm4.3$ \ms\ in L1512, and $21.0\pm5.1$ \ms\ in L1517B \cite{L10a}. We note that similar shifts are indicated for the central parts of L1498 and L1517B in Fig.~11 in \cite{Taf04}, where these targets were observed with the 30-m IRAM telescope also using an \nnhp\ rest frequency which is very close to the CDMS value. Obviously, for more definite conclusions, new laboratory measurements of the rest frequencies and new observations involving other targets and other molecular transitions with different sensitivity coefficients $Q$ are required. It has already been suggested to measure $\Lambda$-doublet lines of the light diatomic molecules OH and CH \cite{Koz09}, microwave inversion-rotational transitions in the partly deuterated ammonia NH$_2$D and ND$_2$H \cite{KLL10}, low-laying rotational transitions in $^{13}$CO and the fine-structure transitions in C\,{\sc i}\ \cite{LMR10}, and tunneling and rotation transitions in the hydronium ion H$_3$O$^+$\ \cite{KL10}. The fourth opportunity is of particular interest since the rest-frame frequencies of H$_3$O$^+$\ transitions are very sensitive to the variation of $\mu$, and their sensitivity coefficients have \textit{different} signs. For example, the two lowest frequency transitions $J_K = 1_1^- \rightarrow 2_1^+$ and $J_K = 3_2^+ \rightarrow 2_2^-$ of para-\hhho\ at, respectively, 307 and 364 GHz have $\Delta Q = Q_{307} - Q_{364} = 14.7$, which is 4 times larger than $\Delta Q = 3.46$ from the ammonia method. This means that the offset $\Delta V \sim 27$ \ms, detected with the ammonia method, should correspond to a relative velocity shift between these transitions, $\Delta V = V_{364} - V_{307}$, of about $100$ \ms\ if $\Delta \mu/\mu \approx 26\times10^{-9}$. We consider the hydronium method \cite{KL10} as an important independent test of the $\Delta \mu/\mu$ value in the Milky Way. To conclude, we note that in cold molecular cores with low ionization degrees ($x_e \sim 10^{-8}-10^{-9}$) frequency shifts caused by external electric and magnetic fields and by the cosmic black body radiation-induced Stark effect are less or about 1 m~s$^{-1}$ and cannot affect the revealed nonzero velocity offset between the rotational and inversion transitions in the ammonia method. Detailed calculations of these effects are given in \cite{L10b}. Our current results tentatively support the hypothesis that the fundamental physical constant~-- the electron-to-proton mass ratio~-- may differ in low-density environments from its terrestrial value. This may be the consequence of the chameleon-like scalar field. However, new laboratory measurements of the molecular rest frequencies and new observations involving other molecular transitions and other targets are required to reach more definite conclusions. \begin{acknowledgement} We are grateful to the staffs of the Medicina 32-m, Nobeyama 45-m, and Effelsberg 100-m radio telescope observatories for excellent support in our observations. We thank Gabriella Schiulaz for her kind assistance in preparing the text. The project has been supported in part by DFG Sonderforschungsbereich SFB 676 Teilprojekt C4, the RFBR grants No. 09-02-12223, 09-02-00352, and 08-02-92001, the Federal Agency for Science and Innovations grant NSh-3769.2010.2, the Program IV.12/2.5 of the Physical Department of the RAS. \end{acknowledgement}

1012.0642

1012

1012.2471_arXiv.txt

{ Although measuring the gas metallicity in galaxies at various redshifts is crucial to constrain galaxy evolutionary scenarios, only rest-frame optical emission lines have been generally used to measure the metallicity. This has prevented us to accurately measure the metallicity of dust-obscured galaxies, and accordingly to understand the chemical evolution of dusty populations, such as ultraluminous infrared galaxies. Here we propose diagnostics of the gas metallicity based on infrared fine structure emission lines, which are nearly unaffected by dust extinction even the most obscured systems. Specifically, we focus on fine-structure lines arising mostly from H{\sc ii} regions, not in photo-dissociation regions, to minimize the dependence and uncertainties of the metallicity diagnostics from various physical parameters. Based on photoionization models, we show that the emission-line flux ratio of ([O{\sc iii}]51.80+[O{\sc iii}]88.33)/[N{\sc iii}]57.21 is an excellent tracer of the gas metallicity. The individual line ratios [O{\sc iii}]51.80/[N{\sc iii}]57.21 or [O{\sc iii}]88.33/[N{\sc iii}]57.21 can also be used as diagnostics of the metallicity, but they suffer a stronger dependence on the gas density. The line ratios [O{\sc iii}]88.33/[O{\sc iii}]51.80 and [N{\sc ii}]121.7/[N{\sc iii}]57.21 can be used to measure and, therefore, account for the dependences on the of the gas density and ionization parameter, respectively. We show that these diagnostic fine-structure lines are detectable with Herschel in luminous infrared galaxies out $z \sim 0.4$. Metallicity measurements with these fine-structure lines will be feasible at relatively high redshift ($z\sim 1$ or more) with SPICA, the future infrared space observatory.

The metallicity of gas and stars in galaxies is one of the most important properties to distinguish various galaxy evolutionary scenarios, since metals result from the cumulative star-formation activity and gas inflow/outflow history in galaxies. Since the gas metallicity of galaxies is in most cases much easier to measure than the stellar metallicity, past studies on the metallicity of galaxies and its evolution have mostly focused on the metallicity of the gas phase (e.g., McCall et al. 1985; Zaritsky et al. 1994; Pettini et al. 2001). Several extensive and detailed studies have been performed on the gas metallicity in local galaxies, the connection with other galaxy properties and the metallicity evolution through the cosmic epochs (e.g. Tremonti et al. 2004; Savaglio et al. 2005; Erb et al. 2006; Maiolino et al. 2008; Mannucci et al. 2009, 2010). Such observational studies have provided strong constraints on evolutionary scenarios (e.g., de Rossi et al. 2007; Kobayashi et al. 2007; Brooks et al. 2007; Finlator \& Dav\'{e} 2008). The vast majority of previous studies on the gas metallicity in galaxies used rest-frame optical diagnostics (see, e.g., Nagao et al. 2006a and references therein). However, such rest-frame optical diagnostics cannot probe the metallicity of the regions affected by significant dust extinction, which is the case for many star forming galaxies, especially at high redshift. This is because the fraction of dusty galaxies such as ultra-luminous infrared galaxies (ULIRGs) increases as a function of redshift (e.g., Le Floc'h et al. 2005). Hence, optical metallicity studies probably probe only the outer, less extinguished region of star forming galaxies, which may deviate substantially from the metallicity of inner more active and more obscured regions of star formation. A clear indication of such a mismatch was recently found in high-z submillimeter galaxies by Santini et al. (2010), who used dust-mass measurements to infer an metals content an order of magnitude higher than inferred from optical metallicity diagnostics. If such mismatch applies to a significant fraction of actively star forming galaxies, this would imply a major revision of the past results on the metallicity evolution in galaxies based on optical diagnostics. Within this context we note that, even for starbursts which have been well-studied in optical, a combination of radio and infrared measurements has conclusively demonstrated that some of the most active star-forming sites are optically obscured (e.g., Gorjian et al. 2001; Vacca et al. 2002; Soifer et al. 2008). The access to far-infrared diagnostics of the gas metallicity would overcome the dust extinction problems plaguing optical measurements. Observational studies on gas metallicity exploiting infrared indicators have already been carried out by using Infrared Space Observatory ($ISO$). Verma et al. (2003) investigated mid-infrared spectra of 12 nearby starburst galaxies and found a strong correlation between gas metallicity and gas excitation, which is also seen in some optical studies (e.g., Nagao et al. 2006; Maier et al. 2006). Garnett et al. (2004) studied chemical properties of gas in M51 based on the $ISO$ data and found that the C/O abundance ratio in M51 is consistent with the solar neighborhood value. This infrared-based metallicity studies is expected to progress dramatically by using the on-going and forthcoming observational facilities such as Herschel, JWST, SPICA, and ALMA. Within this context, it is important to investigate in detail model predictions yielding calibrations between the gas metallicity and flux ratios of emission lines at long wavelengths. Motivated by this, we present the results of photoionization model calculations and the theoretical calibrations of gas metallicity diagnostics based on fine-structure emission lines in mid- and far-infrared wavelength ranges.

\subsection{Comparison with observational data} Here we compare our calculation results on long-wavelength ($\lambda > 50 \mu$m) diagnostic fine structure lines with existing observational data. The only previous instrument that measured the diagnostic fine-structure lines in star-forming galaxies is LWS onboard $ISO$. Specifically we focus on two galaxies, M82 and the Antennae galaxy, since many infrared fine-structure lines are measured for them (Fischer et al. 1996; Colbert et al. 1999). The measured fluxes are summarized in Table 5. Interestingly, the measured flux ratio of [O{\sc iii}]51.80/[O{\sc iii}]88.33 is quite similar between those two galaxies; 1.2 and 1.0 for M82 and the Antennae. These flux ratios suggest $n_{\rm H} \sim 10^2$ cm$^{-3}$ for $U \la 10^{-2}$ (Figure 5). The metallicity diagnostic flux ratio of ([O{\sc iii}]51.80+[O{\sc iii}]88.33)/[N{\sc iii}]57.21 is also similar between M82 and the Antennae, 5.6 and 6.0 respectively. Taking the inferred gas density also into account, these flux ratios correspond to $Z \sim 0.5-0.7 Z_\odot$, i.e., slightly sub-solar metallicities. The observed flux ratio of [N{\sc ii}]121.7/[N{\sc iii}]57.21 in M82 is measured to be 0.5, which, combined with the our initial metallicity estimation, gives $\rm \log U = -3.5$. Only an upper limit of [N{\sc ii}]121.7/[N{\sc iii}]57.21$< 0.28$ (3 $\sigma$) is reported for the Antennae galaxy, providing an upper limit on $\rm \log U$ of about $-3$. Although the constraint on the ionization parameter is not strong, the uncertainty on the ionization parameter does not affect significantly the metallicity estimation. Although both M82 and the Antennae are well-studied nearby galaxies, it is not so straightforward to compare the inferred gas metallicity based on the fine-structure emission-line flux ratios with that inferred from other metallicity diagnostics, because the aperture size of LWS (the beam-size is $\sim 80$ arcsec in FWHM) is very large and difficult to compare with other measurements. Origlia et al. (2004) measured the nuclear gaseous and stellar metallicities by using X-ray and near-infrared spectroscopic observations, and they found that both the metallicities are close to or slightly less than the solar metallicity in M82. Taking the metallicity gradient into account, these metallicities (based also on diagnostics little affected by dust extinction) are consistent with the metallicity inferred through the far-IR fine-structure diagnostics. The situations is similar also for the Antennae galaxy, since its gaseous and stellar metallicities are close to or slightly less than the solar metallicity (Bastian et al. 2006, 2009). Although the comparisons discussed above do not provide a tight test for our new metallicity diagnostics based on far-IR fine structure lines, do at least suggest a broad consistency with other metallicity diagnostics. \subsection{Observational feasibility with next-generation instruments} \begin{figure} \centering \rotatebox{0}{\includegraphics[width=9.6cm]{fig09.eps}} \caption{ Detectability with PACS on Herschel of some diagnostic far-IR fine-structure lines discussed in this paper, which can be used to measure the gas metallicity, ionization parameter and density in galaxies, as a function of redshift. The left hand scale gives minimum $L_{\rm FIR}$ to achieve a 5$\sigma$ detection of each line with an exposure time of 3 hours. The right hand axis shows the minimum star formation rate for the line detectability corresponding to the minimum $L_{\rm FIR}$ based on the relation given in Kennicutt (1999). Green, red, blue, and magenta lines denote [N{\sc ii}]121.7, [N{\sc iii}]57.21, [O{\sc iii}]51.80, and [O{\sc iii}]88.33, respectively. } \label{fig09} \rotatebox{0}{\includegraphics[width=9.6cm]{fig10.eps}} \caption{ Same as Figure 9 but for the detectability with SAFARI boarded on SPICA. } \label{fig10} \end{figure} We finally examine the detectability of the diagnostic fine-structure emission lines discussed here with the new-generation instruments, PACS on Herschel (Poglitsch et al. 2010) and SAFARI (Swinyard 2008; Swinyard et al. 2009) the spectrometer planned for SPICA (a 3m-class cooled telescope in space, Nakagawa 2009). Figure 9 shows the detectability of [N{\sc ii}]121.7, [N{\sc iii}]57.21, [O{\sc iii}]51.80, and [O{\sc iii}]88.33 with PACS. Specifically the figure shows the minimum far-infrared luminosity of galaxies to have each far-IR line detected (S/N=5) with a 3 hour exposure, as a function of redshift. Here we make the simplifying assumption that the luminosity of these emission lines scales roughly with far-infrared luminosity and keeping a ratio similar, on average, to what observed in some nearby galaxies where all of these transitions were observed with past ISO observations; more specifically, $L_{\rm [OIII]88.33}/L_{\rm FIR} \sim 1.8 \times 10^{-3}$, $L_{\rm [OIII]51.80}/L_{\rm FIR} \sim 1.8 \times 10^{-3}$, $L_{\rm [NIII]57.21}/L_{\rm FIR} \sim 6.0 \times 10^{-4}$, and $L_{\rm [NII]121.7}/L_{\rm FIR} \sim 2.7 \times 10^{-4}$ (Fischer et al. 1996; Colbert et al. 1999). The scale on the right hand side gives the minimum SFR corresponding the minimum $\rm L_{FIR}$, by using the SFR--$\rm L_{FIR}$ relationship given in Kennicutt (1999). Figure 9 suggests that we can detect all the fine-structure emission lines needed to infer the metallicity, in luminous infrared galaxies ($\rm L_{FIR}\sim 10^{11}-10^{12} L_{\odot}$) at redshifts up to $z \sim 0.4$. The detection to somewhat higher redshift is restricted to ULIRGs ($\rm L_{FIR}> 10^{12} ~L_{\odot}$). On the other hand, SAFARI boarded on SPICA will drastically expand the detectability, as shown in Figure 10. By using SAFARI, the diagnostic fine-structure lines will be detectable for star-forming galaxies at $z \sim 1$ or at even higher redshift. In Figure 10 we adopt the SAFARI sensitivity given by Swinyard (2008) with a scaling of sensitivity for a conservative 3.0m aperture. Actually the redshift limit may be limited not by the instrument sensitivity, but by the wavelength coverage, although the instrument design of SAFARI has not yet completely fixed.

1012.2471

1012

1012.0474_arXiv.txt

It is argued that the discovery of cosmic acceleration could have been anticipated on thermodynamic grounds, namely, the generalized second law and the approach to equilibrium at large scale factor. Therefore, the existence of dark energy -or equivalently, some modified gravity theory- should have been expected. In general, cosmological models that satisfy the above criteria show compatibility with observational data.

The standard cold dark matter (SCDM) model \cite{peebles} was in good health until around the last decade of the previous century when it became apparent that the fractional density of matter falls well below the Einstein-de Sitter value, $\Omega_{m} = 1$ -see e.g. \cite{mnrs-maddox,nature-efstathiou}. The death blow came at the close of the century with the discovery of the current cosmic acceleration \cite{riess}, something the said model cannot accommodate by any means. However, to account for the acceleration in homogeneous and isotropic models one must either introduce some exotic energy component with a huge negative pressure (dubbed dark energy) or, more drastically, devise some theory of gravity more general than Einstein relativity \cite{reviews}. Thus, both solutions appear somewhat forced and not very aesthetical. Here we argue that dark energy (or something equivalent) is demanded on thermodynamic grounds. In other words, we provide what we believe is a sound thermodynamic motivation for the existence of dark energy. Our argument is based on that the natural tendency of systems to evolve toward thermodynamical equilibrium is characterized by two properties of its entropy function, $S(x)$, namely, it never decreases, $dS(x)/dx \geq 0$, and is convex, $d^{2}S(x)/dx^{2} < 0$ \cite{callen}. In the context of an ever expanding Friedmann-Roberson-Walker (FRW) cosmology this translates in that the entropy of the apparent horizon plus that of matter and fields enclosed by it must fulfill $S'(a) \geq 0$ at any $a$ -the generalized second law (GSL)- as well as $S''(a) \leq 0$ as $a \rightarrow \infty$, where $a$ is the scale factor of the FRW metric and the prime means $d/da$. The apparent horizon in FRW universes always exists (which is not generally true for the particle horizon and the future event horizon) and is known to posses not only an entropy proportional to its area \cite{bak-rey,cai-2008} but also a temperature \cite{cai-2009}. Besides, it appears to be the appropriate thermodynamic boundary \cite{wang-2006}. Before we proceed, it is fair to recall the (at least theoretical) existence of systems lacking any global maximum entropy state, such as Antonov's sphere \cite{lyndenbell68,lyndenbell99}. The latter system consists in a sphere enclosing a number of particles that share some total energy. If the sphere radius happens to increase beyond some critical value, the system becomes unstable and the entropy function ceases to have a global maximum. We shall apply our argument under the assumption that the Universe tends to a state of maximum entropy irrespective of whether it may eventually reach it or not. It would be odd and frustrating that such a universal principle as the second law of thermodynamics could not be applied to the Universe as a whole, especially given the close connection between thermodynamics and gravity \cite{Ted,Pad}. Section II illustrates why dark energy (or some or other modified gravity model) is required on thermodynamic basis and study some dark energy models to see whether they fulfill the thermodynamic criteria. Section III applies the said criteria to some representative modified gravity models. Finally, section IV summarizes and discusses our findings. As is customary, a naught subscript stands for the present value of the corresponding quantity.

As we have argued, neither a radiation nor a cold matter dominated universe can tend to thermodynamic equilibrium in the long run. By contrast, dark energy dominated universes may; this holds true irrespective of whether the dark energy component has entropy or not. Accordingly, dark energy (or something dynamically equivalent at the background level, such as a suitably modified gravity theory) appears thermodynamically motivated. In other words, any of these two kind of ingredients was to be expected on thermodynamic grounds. We, therefore, should not wonder that the Universe is accelerating. However, it does not mean that every accelerating universe is thermodynamically motivated; that is the case, for instance, of any phantom dominated expansion with $w_{x} = {\rm constant}$, and some modified gravity models. One may object that if our reasoning were valid, the Universe would have never ceased to inflate as it would mean a transition from acceleration, ${\cal A}'' < 0$, to deceleration ${\cal A}'' > 0$. Clearly if the primordial inflation would have lasted for ever, big bang nucleo-synthesis would never have occurred, galaxies couldn't have come into existence, and so on -something in stark contrast with observation. However, this reasoning is rather incomplete as it leaves aside the huge entropy generated during the reheating process at the end of inflation. In this explosive and quasi-instantaneous event the inflaton field relinquishes all its energy in the form of matter and radiation and, as a consequence, the Universe sees its temperature enormously increased \cite{lyth-liddle}. Since this huge amount of matter and radiation thermalizes (a necessary condition for primordial nucleosynthesis) both second derivatives, $S''_{m}$ and $S''_{r}$, are negative and, as a result, we may well have that $S''_{r} + S''_{m} + S''_{A} < 0$. Obviously, this will depend on the specific inflationary model and the particular reheating process involved, but we are not aware of any general argument against the fulfillment of this inequality. Specific calculations in this connection will be the subject of a future research. Interestingly enough, cosmological models complying both with the GSL and the thermodynamic criterion that, at late times, the Universe should approach thermodynamic equilibrium appear to be compatible with observational data. Nevertheless, some models that do not comply with the said criteria (as some phantom models and some modified gravity models) seem also consistent with the said data. Our argument could be falsified if it were discovered that in the past (but after the radiation-dominated period set in) the Universe experienced one (or more) transitions from deceleration to acceleration and back. Cosmologies of the type have been proposed to explain the seemingly periodic distribution of galaxies with redshift -see, e.g. \cite{apjl-morikawa}, \cite{grg-nds}- and as an expedient to solve the cosmic coincidence problem \cite{ruth}. However, recent studies on the impact of hypothetical transient periods of acceleration-deceleration on the matter growth \cite{prd_linder} and on the radiation power spectrum \cite{jcap_linder} from the decoupling era, $a \simeq 10^{-5}$, to $a = 0.5$ practically discard these periods in the redshift intervals considered. If eventually it gets confirmed that the present phase of acceleration is to last forever, it may be seen as an indication that the Universe as a whole obeys the laws of thermodynamics (with the reservation that not all models that accelerate at late times comply with them). If, on the contrary, the Universe resumed a decelerated stage one should either call into question the validity of applying the said laws in a cosmological setting or wait for a later, and definitive, accelerating era. Altogether, it is for the reader to decide which possibility looks the less queer: dark energy (or modified gravity) or a universe that will never approach thermodynamic equilibrium. The remaining possibility, either dark energy or modified gravity combined with an increasing departure from equilibrium, involves two oddities rather than one.

1012.0474

1012

1012.2701_arXiv.txt

{ MESS (Mass-loss of Evolved StarS) is a Guaranteed Time Key Program that uses the PACS and SPIRE instruments on board the \it Herschel Space Observatory \rm to observe a representative sample of evolved stars, that include asymptotic giant branch (AGB) and post-AGB stars, planetary nebulae and red supergiants, as well as luminous blue variables, Wolf-Rayet stars and supernova remnants. In total, of order 150 objects are observed in imaging and about 50 objects in spectroscopy. This paper describes the target selection and target list, and the observing strategy. Key science projects are described, and illustrated using results obtained during \it Herschel\rm's science demonstration phase. Aperture photometry is given for the 70 AGB and post-AGB stars observed up to October 17, 2010, which constitutes the largest single uniform database of far-IR and sub-mm fluxes for late-type stars. }

Mass-loss is {\it the} dominating factor in the post-main sequence evolution of almost all stars. For low- and intermediate mass stars (initial mass \less 8 \msol) this takes place mainly on the thermally-pulsing AGB (asymptotic giant branch) in a slow (typically 5-25 \ks) dust driven wind with large mass loss rates (up to 10$^{-4}$ \msolyr, see the contributions in the book edited by Habing \& Olofsson 2003), which is also the driving mechanism for the slightly more massive stars in the Red Supergiant (RSG) phase, while for massive stars (initial mass \more 15 \msol) the mass loss takes place in a fast (hundreds to a few thousand \ks) wind driven by radiation pressure on lines at a moderate rate of a few 10$^{-6}$ \msolyr\ (Puls et al. 2008). Although mass loss is such an important process and has been studied since the late 1960's with the advent of infrared astronomy, many basic questions remain unanswered even after important missions such as {\it IRAS} (Neugebauer et al. 1984), {\it ISO} (Kessler et al. 1996), {\it Spitzer} (Werner et al. 2004) and {\it AKARI} (Murakami et al. 2007): what is the time evolution of the mass-loss rate, what is the geometry of the mass-loss process and how does this influence the shaping of the nebulae seen around the central stars of Planetary Nebulae (PNe) and Luminous Blue Variables (LBVs), can we understand the interaction of these winds with the interstellar medium (ISM) as initially seen by {\it IRAS} (e.g. Stencel et al. 1988) and confirmed by {\it AKARI} (Ueta et al. 2006, 2008) and {\it Spitzer} (Wareing et al. 2006), what kind of dust species are formed at exactly what location in the wind, what are the physical and chemical processes involved in driving the mass-loss itself and how do they depend on the chemical composition of the photospheres? With its improved spatial resolution compared to {\it ISO} and {\it Spitzer}, larger field-of-view, better sensitivity, the extension to longer and unexplored wavelength regions, and medium resolution spectrometers, the combination of the Photodetector Array Camera and Spectrometer (PACS, Poglitsch et al. 2010) and the Spectral and Photometric Imaging Receiver (SPIRE, Griffin et al. 2010) observations on board the {\it Herschel Space Observatory} (Pilbratt et al. 2010) have the potential to lead to a significant improvement in our understanding of the mass-loss phenomenon. This is not only important for a more complete understanding of these evolutionary phases {\it per se}, but has potentially important implications for our understanding of the life cycle of dust and gas in the universe. Dust is not only present and directly observable in our Galaxy and nearby systems like the Magellanic Clouds, but is already abundantly present at very early times in the universe, e.g. in damped Lyman-alpha systems (Pettini et al. 1994), sub-millimetre selected galaxies (Smail et al. 1997) and high-redshift quasars (e.g. Omont et al. 2001, Isaak et al. 2002). The inferred far-IR (FIR) luminosities of samples of $5 < z < 6.4$ quasars are consistent with thermal emission from warm dust ($T < 100$ K), with dust masses in excess of $10^8$ solar masses (Bertoldi et al. 2003, Leipski et al. 2010). It has been typically believed that this dust must have been produced by core-collapse (CC) SuperNovae (SNe), as AGB stellar lifetimes ($10^8$ to $10^9$ yr) are comparable to the age of universe at redshift $>$~6 (Morgan \& Edmunds 2003, Dwek, Galliano \& Jones 2007). The observed mid-IR emission for a limited number of extra-galactic SNe implies dust masses which are generally smaller than $10^{-2}~\rm M_{\odot}$ (e.g. Sugerman et al.\ 2006, Meikle et al.\ 2007, Blair et al. 2007, Rho et al 2008, Wesson et al. 2010a), corresponding to condensation efficiencies which are at least two orders of magnitude smaller than theoretical models predict (Todini \& Ferrara 2001, Bianchi \& Schneider 2007). FIR and sub-mm observations of dust within supernova remnants (SNR) estimate masses ranging from $0.1-1~\rm M_{\odot}$ (Dunne et al. 2003, 2009, Morgan et al. 2003, Gomez et al. 2009), yet there are a number of difficulties with the interpretation of these results. It is obvious that there is now indeed clear observational evidence for dust formation in CCSNe, but the quantity of dust {\it formed} within the ejecta is still a subject of debate. Valiante et al. (2009) recently showed that AGB stars could potentially rival or surpass SNe as the main producer of dust at characteristic timescales of between 150 and 500 Myr, although the model requires rather extreme star formation histories, a top-heavy initial mass function and efficient condensation of dust grains in stellar atmospheres. The dust production of SNe, either from the progenitors (LBV, RSG, Wolf-Rayet (WR) stars) or directly in the ejecta, versus that of AGB stars is therefore of utmost importance and one of the science themes that will be addressed in the MESS Herschel key program described in this paper. \bigskip Most of the astronomical solid state features are found in the near-IR (NIR) and mid-IR (MIR) ranges. The {\it ISO} SWS and LWS spectrometers revolutionised our knowledge of dust and ice around stars. In the LWS range, partly overlapping with Herschel PACS, most of {\it ISO}s spectroscopic dust observations suffered from signal-to-noise (S/N) problems for all but the brightest AGB stars. The sensitivity of Herschel is a clear improvement over {\it ISO} but the short wavelength limit of PACS ($\sim$60\,$\mu$m) is somewhat of a limitation. Nevertheless dust-species like Forsterite (Mg$_2$SiO$_4$) at 69\,$\mu$m, Calcite CaCO$_3$ at 92.6\,$\mu$m, Crystalline water-ice at 61\,$\mu$m, and Hibonite CaAl$_{12}$O$_{19}$ at 78\,$\mu$m are expected to be detected. Other measured features lack an identification e.g. the 62-63\,$\mu$m feature with candidate substances like Dolomite, Ankerite, or Diopside (see Waters 2004 and Henning 2010 for overviews). At longer wavelengths, PAH `drum-head' or `flopping modes' have been predicted to occur (Joblin et al. 2002), that can be looked for with the SPIRE FTS (Fourier Transform Spectrometer) that will observe in an previously unexplored wavelength regime. Apart from solid state features the PACS and SPIRE range contain a wealth of molecular lines. Depending on chemistry and excitation requirements, the different molecules sample the conditions in different parts of a circumstellar envelope (CSE). While for example CO observations in the J= 7-6 line (370~$\mu$m) can be obtained under good weather conditions from the ground, this line traces gas of about 100~K. With SPIRE and PACS one can detect CO J= 45-44 at 58.5~$\mu$m at the short wavelength edge of PACS (as was demonstrated in Decin et al. 2010a) which probe regions very close to the star. Although only the Heterodyne Instrument for the Far Infrared (HIFI, de Graauw et al. 2010) onboard \it Herschel \rm will deliver resolved spectral line observations, PACS and SPIRE with their high throughput will allow full spectral inventories to be made. The analysis of PACS, SPIRE (and HIFI and ground-based) molecular line data with sophisticated radiative transfer codes (e.g. Morris et al. 1985, Groenewegen 1994, Decin et al. 2006, 2007) will allow quantitative statements about molecular abundances, the velocity structure in the acceleration zone close to the star, and (variations in) the mass-loss rate. \medskip With these science themes in mind, the preparation for a Guaranteed Time (GT) Key Program (KP) started in 2003, culminating in the submission and acceptance of the MESS (Mass-loss of Evolved StarS) GTKP in June 2007. It involves PACS GT holders from Belgium, Austria and Germany, the SPIRE {\it Specialist Astronomy Group} 6, and contributions from the {\it Herschel Science Centre}, and Mission Scientists. The allocated time is about 300 hours, of which 170h are devoted to imaging and the remaining to spectroscopy. \medskip Section~2 describes the selection of the targets and Section~3 describes the observing strategy. Section~4 discusses some aspects of the current data reduction strategy. Section~5 presents the key science topics that will be pursued and this is illustrated by highlights of the results obtained in the Science Demonstration Phase (SDP), and presenting ongoing efforts. Aperture photometry for 70 AGB and post-AGB stars is presented and compared to {\it AKARI} data. Section~6 concludes this paper. In two appendices details on the PACS mapping and data reduction strategy are presented.

The scope, aims and status of the \it Herschel \rm Guaranteed Time Key Program MESS (Mass-loss of Evolved StarS) are presented. The current concepts on the data processing are presented, and aperture photometry of all 70 AGB and post-AGB stars observed as per October 17th 2010 are presented\footnote{The progress of the MESS project can be followed via www.univie.ac.at/space/MESS}. Some of our SDP data is already public (see the remarks in Table~\ref{list-all}) and the data taken in routine phase have a proprietary period of one year, implying that data will become successively public from about December 2010 onwards. Currently, the PACS maps are produced via the standard {\tt PhotProject} task, that use data that are filtered to remove 1/f-noise as input. We achieve good results for compact objects, but unmasked large-scale structures in the background are affected by the filtering. In order to improve this, we are currently investigating other filtering and mapping methods, such as MADmap (Cantalupo et al. 2010). That way we also seek to improve the spatial resolution of the final maps. A second point under investigation is to correct the effects caused by the instrument PSF. With its tri-lobe pattern and other wide-stretched features it is currently not possible to make any definite statements about structure in the circumstellar emission close to the central object, although many sources are extended. Thus we are investigating different deconvolution strategies and PSF-related matters. The efforts the MESS consortium are currently undertaking to improve the PACS data processing are described in Ottensamer et al. (2011). On the science side, the publication of the very first {\it Nature} paper based on {\it Herschel} results by Decin et al. (2010b) is a highlight. It demonstrates the power of the PACS and SPIRE spectrometers. With more than 50 PACS and almost 30 SPIRE targets to be observed spectroscopically this will result in an extremely rich database that, with proper modelling, will allow detailed studies on molecular abundances, the velocity structure in the acceleration zone close to the star, and the mass-loss rate. On the imaging side the fact that bow shock cases are ubiquitous is extremely interesting. Although this in fact makes it more difficult to derive the mass-loss rate history of the AGB star, it offers an unique opportunity to use these cases as probes of the ISM.

1012.2701

1012

1012.0532_arXiv.txt

{Recently, it has been demonstrated that neutrinos in a supernova oscillate \emph{collectively}. This process occurs much deeper than the conventional matter-induced MSW effect and hence may have an impact on nucleosynthesis. In this paper we explore the effects of collective neutrino oscillations on the $r$-process, using representative late-time neutrino spectra and outflow models. We find that accurate modeling of the collective oscillations is essential for this analysis. As an illustration, the often-used ``single-angle'' approximation makes grossly inaccurate predictions for the yields in our setup. With the proper multiangle treatment, the effect of the oscillations is found to be less dramatic, but still significant. Since the oscillation patterns are sensitive to the details of the emitted fluxes and the sign of the neutrino mass hierarchy, so are the $r$-process yields. The magnitude of the effect also depends sensitively on the astrophysical conditions --- in particular on the interplay between the time when nuclei begin to exist in significant numbers and the time when the collective oscillation begins. A more definitive understanding of the astrophysical conditions, and accurate modeling of the collective oscillations for those conditions, is necessary. }

There are a number of types of element synthesis that are thought to occur in ``hot outflows'', such as the formation of light $p$-process elements, possibly $r$-process elements, and also some iron peak elements. A hot outflow is one in which nucleons are originally photo-dissociated into free nucleons at temperatures of around a few tens of MeV. Material typically flows away from a hot center, such as an accretion disk around a black hole or proto-neutron star, and as it does so the nucleons cool and combine into nuclei. Given the high temperature of the center, neutrinos can be emitted in vast numbers. These neutrinos interact with the outflowing and cooling nucleons and nuclei and, among other things, change neutrons into protons and vice-versa. Thus any process which alters the luminosity or spectra of the neutrinos may impact element synthesis. Neutrino flavor transformations -- and, specifically, collective oscillations, as explained later -- are a prime example of such a process and it is essential to determine whether they are important for nucleosynthesis in hot outflow environments. In this initial work we explore, by considering a few specific examples, the size of this effect on the abundances of rapid neutron capture elements. It should be stressed that the physics underlying the oscillations is by now well established and treating it is no longer optional. The essential ingredients are the measured neutrino oscillation parameters and, crucially, the neutrino-neutrino coherent interactions \cite{Fuller:1987,Notzold:1987ik,Pantaleone:1992xh,Pantaleone:1992eq,Sigl:1992fn,McKellar:1992ja}, which can cause \emph{collective} oscillations. It has been suspected for nearly two decades that neutrino oscillations ({\it e.g.}, \cite{Qian:1993dg}), particularly the collective oscillations between the active flavors \cite{Pantaleone:1994ns,Qian:1994wh,Sigl:1994hc,Pastor:2002we,Balantekin:2004ug}, could potentially impact the $r$-process. The case for this became much stronger in the last five years, since supercomputer calculations of collective oscillations have become available (starting with \cite{Duan:2006an,Duan:2006jv}). These calculations have shown that, even without sterile neutrinos, large flavor transformations can develop sufficiently close to the neutrinosphere, where the $r$-process is thought to occur. Yet, with the exception of the recent attempt \cite{Chakraborty:2009ej}, the impact of the new results on the $r$-process has not been investigated. The reason for this, we believe, is not the lack of interest, but the inherent complexity of the problem. First, accurate modeling of collective oscillations requires supercomputers. This is so because -- unlike the more familiar MSW flavor transformations, in which each neutrino evolves independently -- in the collective oscillation regime neutrinos of different directions and different energies are coupled \cite{Qian:1994wh,Duan:2006an}. Second, one has a chain of nuclear reactions to follow, as a function of the position in the outflow. Lastly, the physical conditions, in which the candidate $r$-process takes place, themselves require detailed simulations. In its full form, this is clearly a very ambitious program. The goal of the present paper is to gauge the effect and establish whether or not its magnitude warrants further, more detailed modeling. Moreover, we would like to get a sense of the size of the effect for a variety of conditions. With these goals in mind, we pick astrophysical conditions which produce neutron rich outflow, following models in the existing literature. Since collective effects may cause rapid oscillations above the neutrinosphere, we couple neutrino transformation calculations to a nuclear reaction network. We present several examples of calculations obtained in this way. One may attempt to reduce the impact of the oscillations on the $r$-process to a simple criterion, such at the electron fraction $Y_{e}$ at a certain radius. Such prescriptions should also be applied with care and we discuss their limitations. One is also tempted to simplify the treatment of the collective oscillations. We show, however, that such simplifications may lead to entirely erroneous results. We illustrate this with the example of the frequently used ``single-angle'' approximation, which in this case is seen to utterly fail. This sensitivity to the details of the oscillations -- and, correspondingly, the need to model this process carefully -- is one of the major results of this paper. Lastly, although our outflow models are clearly simplified (and, indeed, as explained in the next Section, the exact mechanism of the astrophysical $r$-process is yet to be conclusively established), we believe our basic framework and findings should apply more generally. For example, NS-NS and BH-NS mergers also have neutrinos which, for hot outflows, strongly influence the composition of the accretion disk and outflow \cite{Surman:2005kf,Met08}. We will return to this important point at the end. This paper is organized as follows. In Sect.~\ref{sect:overview}, we briefly summarize the current status of the $r$-process modeling. In Sect.~\ref{sect:toy}, we illustrate the effects of neutrino flavor transformations using a toy model. In Sect.~\ref{sec:collective}, we introduce the main physics of the collective transformations and give order-of-magnitude arguments why these transformations may impact the $r$-process. In Sect.~\ref{sec:nuc_calc}, we describe the nucleosynthesis network calculations and, in Sect.~\ref{sec:collectiverprocess} we present our sample results. Finally, in Sect.~\ref{sec:conclusions}, we summarize our findings and discuss them in a more general context of astrophysical nucleosynthesis.

\label{sec:conclusions} We establish, using full three neutrino flavor multi-angle calculations, that the flavor transformation can begin in regions sufficiently close to the center of a proto-neutron star that the outcome of element synthesis in hot outflows is appreciably affected. Further, the effect on the final abundance distribution is comparable to that which stems from uncertainties in the nuclear conditions, e.g. \cite{Surman:2008ef,Arcones:2010dz}. Neutrino collective oscillations act to enhance the role that neutrinos play in $r$-process nucleosynthesis, and decreases the efficacy of the rapid neutron capture process. We find that the effect of collective flavor transformation can be dramatically overestimated if the ``single-angle'' approximation is made. An observable effect on the abundance pattern can occur in both the normal and inverted hierarchy. We find that the outcome is dependent not only on the spectra and luminosity of the electron and antielectron neutrinos but also on the $\mu$- and $\tau$-type neutrino spectra. We stress again that the nucleosynthesis calculations are sensitive to how collective oscillations develop early on, whereas neutrino signal in a terrestrial detector measures the final spectra after the oscillations cease. There are known cases where the final spectra come out qualitatively similar in the single-angle and multiangle calculations, even though at the intermediate stages the spectra are radically different \cite{Duan:2010bf}. In this sense, the nucleosynthesis calculations place high demands on the accuracy of the oscillations calculations. We continue the tradition of studying neutron rich outflows from proto-neutron stars that might occur in supernovae. While self-consistent models which produce a sufficiently neutron rich environment in these conditions have remained elusive, this site remains under consideration as a potential environment due to a variety of astrophysical indicators that favor core-collapse supernovae as the origin of the ``main'' $r$-process elements. If and when a self consistent model of a hot outflow is produced, we have shown that these neutrino flavor transformation effects will need to be included. In the scenarios we have considered, the flavor transformation effects make it more difficult to produce $r$-process nuclei. However, one must keep in mind that if new, more neutron rich possibilities are discovered, the flavor transformation effects could potentially work in the direction of improving the comparison with abundance data. Our work improves on previous efforts to estimate the effect of the $r$-process in two important ways. (1) We use three flavor multi-angle calculations as opposed to ``single angle'' and/or two flavor calculations. (2) We couple our three flavor multi-angle calculation to a nuclear reaction network so that we include the effects of material composition. These are both essential for making reasonable quantitative estimates of the impact of flavor transformation. There are a variety of types of elements that are made in the presence of a strong neutrino flux. This type of environment occurs in supernovae near proto-neutron stars, and in compact object mergers and gamma ray bursts from disks around black holes, e.g.~\cite{Surman:2005kf,Surman:2008qf,Metzger:2007}. In these strong neutrino fluxes, as we have discussed, neutrinos can be expected to transform collectively. While it may be possible to find some environments with hot outflows where the abundance yields are unaffected by neutrino flavor transformation, this is not clear a priori. The analysis presented here suggests that flavor transformation must be considered in order to predict accurate abundance yields in environments with strong neutrino fluxes.

1012.0532

1012

1012.5121_arXiv.txt

{ It is well known that global symmetries protect local supersymmetry and a zero value for the cosmological constant in no--scale supergravity. The breakdown of these symmetries, which ensure the vanishing of the vacuum energy density, results in a set of degenerate vacua with broken and unbroken supersymmetry leading to the natural realisation of the multiple point principle (MPP). Assuming the degeneracy of vacua with broken and unbroken SUSY in the hidden sector we estimate the value of the cosmological constant. We argue that the observed value of the dark energy density can be reproduced in the split-SUSY scenario if the SUSY breaking scale is of the order of $10^{10}\,\mbox{GeV}$.} \FullConference{35th International Conference of High Energy Physics\\ July 22-28, 2010\\ Paris, France} \begin{document}

1012.5121

1012

1012.5341_arXiv.txt

Rankine-Hugoniot condition has been solved to study phase transition in astrophysical scenario mainly in the case of phase transition from neutron star (NS) to quark star (QS). The phase transition is brought about by a combustion front, which travels from the center to the surface. The equations of state and temperature plays a huge role in determining the nature of the front propagation, which brings about the phase transition in neutron stars (NSs). Magnetic field has been introduced and the modified conservation condition for the perpendicular and oblique shocks is obtained. Numerical solution of the perpendicular shock has been shown in the figures, which finds that the magnetic field helps in shock generation. It indirectly hints at the instability of the matter and thereby the NS for very high magnetic field, implying that NSs can only support finite magnetic field strength.

When the velocity of a fluid in motion becomes comparable with or exceeds that of the sound, the effect due to compressibility of the fluid become of prime importance. For a wave propagating in a non conducting gas, when the amplitude is so small that linear theory applies, the disturbance propagates as a sound wave. If the gas has a uniform pressure and density, the speed of propagation of sound and the wave profile maintains a fixed shape, since each part of the wave moves with same speed. However, when the wave possesses a finite amplitude, so that nonlinear terms in the equation become important, the crest of the sound wave moves faster than its leading or trailing edge. This causes a progressive steepening of the front portion of the wave as the crest catches up and ultimately, the gradient of pressure, density, temperature and velocity become large that dissipative processes, such as viscosity or thermal conduction are no longer negligible. Then a steady wave shape is attained, called a shock wave. The shock wave moves at a speed in excess of the sound speed, so the information cannot be propagated ahead to signal its imminent arrival, since such information would travel at sound speed relative to the undisturbed medium ahead of the shock. The dissipation inside the shock front leads to a gradual conversion of the energy being carried by the wave into heat. Thus, the effect of the passage of a shock wave are to convert ordered (flow) energy into random (thermal) energy through particle collisions and also to compress and heat the gas. The shock front itself is in reality a very thin transition region. Its width is typically only a few mean-free paths, with particle collisions establishing the new uniform state behind the shock. The relativistic shock propagates into a medium with a changing equation of state. Therefore, a simple analysis of the jump condition for a polytropic or perfect fluid is not adequate and a deep understanding of this problem calls on for the full theoretical description of the relativistic shock in a medium with arbitrary equation of state. Further complication might arise if there is a presence of significant magnetic field. In fact, one can show that the relative importance of a magnetic field can grow during a collapse. In the field of nuclear physics, high energy collisions among heavy ions can be modeled by using fluid dynamical concepts. Also, some current models under investigation predict that relativistic shocks (or relativistic detonation and deflagration) might be related to the phase transition from nuclear matter to a quark matter. Relativistic shock waves have been the subject of early investigation in relativistic fluid dynamics and magneto-fluid dynamics. In relativistic fluid dynamics the pioneering work is that of Taub \cite{taub} where the relativistic form of jump condition is established. A detailed analysis of the thermodynamic treatment of classical shock waves, is due to Thorne \cite{thorne}. Explicit solutions of the jump condition have ben obtained for special equation of states. Shock wave in relativistic magneto-fluid have been investigated extensively and in rigorous mathematical way by Lichnerowicz. \cite{lichnero}. Detonation and deflagration waves in relativistic magneto-fluid dynamics for nuclear physics and cosmology have been investigated by \cite{stein,cleymans}. In Astrophysical scenario shock plays very important role in determining the outcome of compact stars. In the case of massive stars, in the range between 8-100 solar masses, which are thought to be the progenitors of type II supernovas, one of the most viable mechanism for producing a explosion is gravitational collapse and bounce \cite{van}. In this case a shock is formed outside the inner core, which propagates outwards reaching relativistic speeds. Shock waves are also responsible for phase transition and gamma ray bursts (GRB) in compact stars. In this paper I will mostly concentrate on the phase transition scenario in a compact star. Witten \cite{key-1} conjectured of strange quark matter (SQM), consisting of approximately equal numbers of up ({\it u}), down ({\it d}) and strange ({\it s}) quarks, is believed to be the ground state of strong interaction. This was supported by model calculations for certain ranges of values for strange quark mass and strong coupling constant \cite{key-2}. After that there has been constant efforts at confirming the existence of and SQM, though transiently, in ultra relativistic collisions. On the other hand, SQM could naturally occur in the cores of compact stars, where central densities are expected to be an order of magnitude higher than the nuclear matter saturation density. Thus, neutron stars which have sufficient high central densities might convert to strange star, or at least hybrid (a star with a quark core) stars. These transitions may lead to various observable signatures in the form of a jump in the breaking index and gamma ray bursts \cite{key-3,key-4}, and a full QS might help in explaining the phenomena of observed quasi periodic oscillations \cite{key-4a}. \par There may be several scenarios by which neutron stars could convert to quark stars. It may happen through a "seed" of external SQM \cite{key-5}, or triggered by the rise in the central density due to a sudden spin-down in older neutron stars \cite{key-6}. Several authors have studied the conversion of nuclear matter to strange matter under different assumptions \cite{key-7,key-8,key-9,key-10,key-11,key-12,key-13, key-14,key-15,key-16,key-17}. They have been summarized in a recent work of ours \cite{key-18} and for the constraint of space, I do not repeat them here. After the discovery magnetars, some compact stars were found to have very high surface magnetic fields. So the study of relativistic rankine-hugoniot condition is not sufficient. To have a full understanding of the properties of NS and its phase transition to QS, such conditions should be examined in the presence of high magnetic fields. In this paper I wish to carry out such a basic calculation keeping our focus mainly on the astrophysical scenario. The paper is organized as followed: first I will discuss the general rankine hugoniot condition as a discontinuity in the conversion front. In section III, I will introduce magnetic field and study the new set of modified conservation equations. Then in section IV, I shall show my results and finally I will discuss and summarize them.

Finally in this section I summarize my results. I find that rankine-hugoniot condition can be solved to determine the condition for different types of wave generation in a neutron star. It also determines the mode of the propagation of the wave front. The temperature (finite temperature of the matter) helps in the generation of the front and as the temperature rises the wave front can generate both at much lower and at much higher baryon densities which was not possible for the zero temperature case. Next I write down the modified conservation conditions in presence of magnetic field. I have written down conditions for both the perpendicular and oblique shock waves. The inclusion of the magnetic field introduces not only extra conditions but also the earlier existing conditions gets modified. I have solved the velocity of the matter of the two phases and plotted curves for the perpendicular wave, which was not obtained analytically. The conditions cannot be solved analytically and and therefore I have solved the nonlinear simultaneously equation numerically. I have matched my nonmagnetic numerical results with the analytically solvable nonmagnetic solutions. The oblique wave equation gets very much complicated and the simultaneous equations does not converges for different value of the baryon densities. So I have not plotted the results of the oblique waves. Solving the perpendicular wave for finite magnetic field I find that the range of baryon density, for which the flow velocities of matter are physical, increases with magnetic field strength. This is because by the introduction of the magnetic field the resultant pressure and energy redistribute in such a way that, for same baryon density, the changes for shock generation increases. I also find that there is a cut off magnetic field strength that matter can support, beyond which the matter becomes unstable and the flow velocities becomes imaginary. This on the other hand suggest that a NS cannot support magnetic field beyond a certain field strength. To finally summarize our result I mention that this is the first instance where such a treatment of modified conservation condition has been done in the presence of magnetic field in the astrophysical phase transition scenario. This provides with new interesting results, that has not been anticipated before and also indirectly hints at the instability of the NS at very high magnetic field. More interesting results is anticipated if the full oblique wave equations can be solved and I am now focussing mainly on that problem. I would like to thank Grant No. SR/S2HEP12/2007, funded by DST, India for financial support.

1012.5341

1012

1012.2082_arXiv.txt

We present methods and preliminary results of a relatively novel search for nearby stars. The method relies on photometric distance estimates as its primary search criterion, thus distinguishing itself from proper motion-based searches that have produced the bulk of nearby star discoveries.

Over the last 200 years, proper motion (the apparent motion of stars across the sky, seen over periods of years) has been used to find nearby stars. This approach has been based on the idea that stars move through space, and the closest ones should appear to move the fastest. Currently, we understand that this is due to the combination of the motions of the Sun and the star in question in their orbits around the center of the Galaxy, though this idea predates the discovery of the Galaxy and our place in it by at least 150 years. The earliest reference available seems to be William Herschel \citep{1783RSPT...73..247H}, who claims the phenomena is well-established and credits its discovery to Sir Edmund Halley. This property of large proper motion has served nearby star research well from the very beginning, forming at least part of the decisions of \citet{1838MNRAS...4..152B} and \citet{1839MNRAS...4..168H} to observe 61 Cygni and Alpha Centauri (respectively) for parallax. The trend continues to the present. Nearly all nearby stars known have high proper motions, with the limit of high proper motion (0.5\arcsec yr$^{-1}$, van Maanen) or interesting proper motion (0.2\arcsec yr$^{-1}$, the Royal Greenwich Observatory) set by influential publications in the early part of the 20th century \citep{1988IAUS..133..301L}. Though Luyten does not give the reasoning behind either argument, we suspect the 0.2\arcsec yr$^{-1}$~limit was likely the best that could be done visually with the photographic plates and blink comparators of the time. Still, there are signs proper motion is not an entirely foolproof method. Proxima Centauri, the closest star to Earth, is only the 18th highest proper motion star in the New Luyten Two Tenths Catalog \citep[][, though the other 17 are also very nearby stars]{1979nltt.book.....L} (though the other 17 are also very nearby stars). By the same token, many brighter stars were singled out for parallax, e.g. Hipparcos observed all stars brighter than $V$=7.3 \citep{1997A&A...323L..49P} and discovered not all objects within 25 parsecs are fast moving. One good example is Gl 566, which despite a distance of 6.7 parsecs is moving at 0.169\arcsec yr$^{-1}$. It would likely not have been noticed except that it is a 5th magnitude G star with a K companion and visible orbital motion. In any case, years later we still have unfinished business. The situation is currently thus: assuming that 50 parallax-verified systems within 5 parsecs is a statistically meaningful sample, there should be 400 stars within 10 parsecs (8 times the volume), and 6250 systems within 25 parsecs (125 times the volume). The current tally is 256 parallax-verified systems within 10 parsecs\footnote{RECONS 10 pc census, http://www.recons.org/census.posted.htm; Henry, T.J. accessed 2010-12-01} and 2011 parallax-verified systems within 25 parsecs \citep[NStars,][]{2002AJ....123.2002H}. We are missing nearly 36\% of systems within 10 pc and 68\% of all stars within 25 pc.

1012.2082

1012

1012.2561_arXiv.txt

Continuum intensity observations obtained with the Michelson Doppler Imager (MDI) on-board the SoHO mission provide long time series of filtergrams that are ideal for studying the evolution of large-scale phenomena in the solar atmosphere and their dependence on solar activity. These filtergrams, however, are not taken in a pure continuum spectral band, but are constructed from a proxy, namely a combination of filtergrams sampling the \ion{Ni}{1} 6768 \AA\ line. We studied the sensitivity of this continuum proxy to the shape of the nickel line and to the degradation in the instrumental transmission profiles. We compared continuum intensity measurements in the nearby of nickel line with MDI proxy values in three sets of high resolution spectro-polarimetric data obtained with the Interferometric Bidimensional Spectrometer (IBIS), and in synthetic data, obtained from multi-dimensional simulations of magneto-convection and one-dimensional atmosphere models. We found that MDI continuum measurements require brightness corrections which depend on magnetic field strength, temperature and, to a smaller extent, plasma velocity. The correction ranges from 2\% to 25\% in sunspots, and is, on average, less than 2\% for other features. The brightness correction also varies with position on the disk, with larger variations obtained for sunspots, and smaller variations obtained for quiet sun, faculae and micropores. Correction factors derived from observations agree with those deduced from the numerical simulations when observational effects are taken into account. Finally, we found that the investigated potential uncertainties in the transmission characteristics of MDI filters only slightly affect the brightness correction to proxy measurements.

The continuum intensity measurements obtained with MDI have provided extensive time series of data unaffected by seeing for over 13 years. These data are ideal for studying large-scale phenomena in the solar atmosphere and their dependence on solar magnetic activity. They have been utilized in many investigations, concerning the analysis of large-scale patterns in plasma motions \citep[e.g.,][and references therein]{meunier2008,meunier2007}, the measurement of the radiative properties of magnetic elements over the activity cycle \citep[e.g.][]{ortiz2002,ortiz2006,mathew2007}, and the modeling of irradiance variations \citep[e.g.][]{krivova2003,wenzler2006,wenzler2009}. Nevertheless, MDI continuum data are a by-product of the instrument, which was designed mainly for Doppler measurements. Specifically, intensities in the continuum are derived by the combination of five narrow-band filtergrams, obtained with filters sampling a passband of 94 m\AA\ FWHM, equally spaced by 75 m\AA\ around the \ion{Ni}{1} 6768 \AA\ mid-photospheric line. The filtergrams are labeled $F_0 \ldots F_4$, where $F_0$ is divided over two bands taken near the continuum, on either side of the line, $F_1$ and $F_4$ are centered on the wings of the nickel line, and $F_2$ and $F_3$ are centered around its core. Even at $F_0$ the instrument does not sample true continuum. Instead, a proxy continuum-intensity filtergram is constructed by combining the five nominal filtergrams in the following way \citep{scherrer1995}: \begin{equation} I_c = 2 F_0 + I_{\rm{depth}}/2 + I_{\rm{ave}}, \label{eq:IcontMDI} \end{equation} where $I_{\rm{ave}}$ is the average of the five filtergrams and $I_{\rm{depth}}$ is the line depth, which is given by: \begin{equation} I_{\rm{depth}} = \sqrt{2((F_1 - F_3)^2 + (F_2 - F_4)^2)}. \label{eq:linedepth} \end{equation} The components of sum $I_c$ theoretically have cancelling systematic errors as a function of solar velocity, so that the continuum image is claimed to be free of Doppler induced cross-talk at the 0.2\% level \citep{scherrer1995}. The accuracy of the MDI continuum estimate (called the MDI method hereafter) is, however, inherently limited by the shape of the spectral line and its sensitivity to both strength and inclination of the magnetic field, local thermal stratification, and line-of-sight motions of the observed region. In addition, MDI measurements may also suffer from substantial errors because of uncertainties in the actual transmission characteristics of the filters utilized \citep{wachter2008}. Several previous studies have focused on the accuracy of magnetic flux estimates obtained from MDI observations \citep[e.g.][]{berger2003,tran2005,ulrich2009,wang2009, demidov2009,zhendong2010}, as well as the accuracy of dynamic measurements \citep[][]{wachter2006,rajaguru2007,wachter2008}. Results indicate a substantial underestimate of magnetic flux and spurious contributions affecting Doppler measurements. This lead to recalibration of magnetograms and dopplergrams series, with the most recent results suggesting that the calibration of these data could be improved yet further. By contrast, the accuracy of continuum intensity measurements has only been investigated by \citet{mathew2007}, using spectral calculations based on one-dimensional atmospheric models. Their results indicate that MDI continuum measurements underestimate intensity in sunspots, with an error depending on both the magnetic field strength and temperature stratification. The objective of this study is to further investigate the uncertainties affecting MDI continuum intensity measurements, by considering both high spatial resolution spectro-polarimetric observations and numerical simulations. To this aim we defined a brightness correction factor and investigated its variation with the physical properties of the analyzed solar features and observational conditions, such as spatial and spectral scattered light, finite spatial and spectral resolution and position on the solar disk. We also investigated possible effects resulting from degradation of MDI filter transmission profiles. The paper is organized as follows. We describe the data analyzed and provide some details of the reduction applied (Sect. 2). Then, we present the spectral synthesis method and the results obtained from our measurements and synthesis computations (Sect. 3). We also investigate sensitivity of results to filter profiles uncertainties (Sect. 4). We finally discuss the results obtained and present our conclusions (Sect. 5).

We have presented a study of the accuracy of the MDI continuum intensity measurement, with the help of high spatial resolution spectro-polarimetric observations acquired with the IBIS/DST and outcomes from numerical simulations representative of different solar regions. In order to provide the reader with a quantity that clearly indicates the MDI uncertainties and allows their compensation on measurements, we defined as brightness correction the ratio between the contrast value in the continuum and the one estimated through MDI method and investigated the variation of this quantity with the physical properties of analyzed solar regions. We found that the MDI brightness correction, on average, increases with the increase of the LOS magnetic field strength, the decrease of temperature, and the increase of LOS velocity. To summarize, the contrast is overestimated in quiet upflow regions, and is underestimated in magnetic and quiet downflow regions. The amount of correction is less than 1\% in quiet regions and increases with the magnetic field strengths, reaching values up to 25\% of the MDI continuum value in sunspots. However, line saturation effects limit the MDI error in regions with highest magnetic field strengths. Results derived from observations agree with those deduced from simulations, when accounting for the observational conditions and instrumental degradations affecting the observations. Among the effects considered, spatially spurious light is the one that affects the most the brightness correction values. In the lack of a reliable estimate on our observational data, we therefore refrained from a direct comparison of results from MDI and IBIS data. On the other hand, we did not find any significant variation of the brightness correction values when degrading the results from IBIS data to the spatial resolution of MDI, thus suggesting that brightness correction is not significantly dependent on the spatial resolution of the observations. Note that the brightness correction values that we found correspond to errors much higher than the 0.2\% quoted in the instrument description of \citet[][p.~163]{scherrer1995}. The brightness correction factor derived from synthetic spectra also depends on the position on the solar disk in different manners for the different features analyzed. We found that the brightness correction for sunspot models with strong field ($B_{\mathrm{LOS}} \ge 3$ kG) varies by up 2\% for disk positions corresponding to $0.9 \leq \mu \leq 1$, reaching a maximum value for $\mu = 0.9$, and then decreasing towards the limb. For weaker fields the variation across the disk is much less (Fig.~\ref{fig_CLV}). Also for quiet Sun, faculae and micropores the variation in brightness correction is small, below 1\%, across the disk. Finally, we investigated effects of possible filter degradations on MDI continuum contrast estimates. We estimated variations with the shift of the passband pattern and with the change in the width of the transmission profiles of the instrument. Results are summarized in Fig. \ref{fig_aber_shift} for various features observed on the solar photosphere. We found that for small amounts ([-0.03,0.03] \AA) of transmission profiles displacement the variation of the brightness correction is negligible for features not affected by Doppler shifts (i.e. faculae and quiet sun). For line profiles shifted with respect to quiet sun line center, a shift of filter center in the same direction as the line Doppler shift causes a reduction of the amount of correction, whereas a displacement in the opposite direction causes an increase. In addition, we found that broadening of the passband causes a decrease of the amount of correction when the width of the line profiles is larger than that of the passband. For transmission line profiles wider than the profile of the solar feature under study, the contrast is overestimated and the amount of deviation of the correction coefficient from unity increases with the increase of the filter degradation. Our results are qualitatively in agreement with those presented by \citet{mathew2007}, who applied MDI procedure to synthetic line profiles obtained from \citet{kurucz91} atmosphere models. Nevertheless, a comparison of results shows that they obtained brightness correction larger than the ones we have presented. These discrepancies must be ascribed to the differences in model atmospheres employed, assumptions for atomic parameters and micro and macro velocity values. Due to the arbitrariness of the choice of these quantities inherent to 1D models, this analyses should be repeated by a direct comparison of MDI data with observations in the \ion{Ni}{1} continuum properly compensated for scattered light, or with the employment of magneto-convection simulations encompassing a larger variety of sunspots. Results presented in this study show that photometric measurements obtained with MDI data should be revised by taking into account the brightness correction factors presented. Due to the dependence of this factor on observational conditions, continuum images should be compensated for spatial scattered light before any correction is applied. Since MDI Point Spread Function varies with time, this result contributes to question studies based on the temporal variation of photometric properties of solar features derived from MDI data analyses. Finally, our study proofs forward modeling as a powerful tool also for instrumental calibration. In particular, results presented here are of interest for the interpretation of data acquired with the Helioseismic and Magnetic Imager onboard the Solar Dynamic Observatory.

1012.2561

1012

1012.4940_arXiv.txt

{% The metal-poor, fundamental-mode (P0) and first-overtone (P1) Cepheids in the dwarf galaxies IC\,1613, WLM, Pegasus, Sextans~A, Sextans~B, and Leo~A are compared with the about equally metal-poor Cepheids of the Small Magellanic Cloud (SMC). The period-color (P-C) and period-luminosity (P-L) relations of the seven galaxies are indistinguishable, but differ distinctly from those in the Large Magellanic Cloud (LMC) and the solar neighborhood. Adopting $(m-M)^{0}_{\rm SMC}=18.93$ from independent evidence, one can determine reliable distance moduli for the other dwarf galaxies of $(m-M)^{0}=24.34\pm0.03$, $24.95\pm0.03$, $24.87\pm0.06$, $25.60\pm0.03$ (mean for Sextans~A \&~B), and $24.59\pm0.03$, respectively. }

\label{sec:1} It has been shown that the character of the period-luminosity (P-L) relations varies particularly at short wavelengths as a function of the metallicity (\citealt*{STR:09}, in the following Paper~III). This implies the {\em prediction\/} that the very low-metallicity Cepheids in IC\,1613, WLM, and the Pegasus dwarf system should follow P-L relations that are more similar to those of the SMC than those defined by the more metal-rich Cepheids in the LMC and the Galaxy. The same prediction holds for the period-color (P-C) relations. The purpose of this paper is therefore to compare the P-L and P-C relations of the three above-mentioned galaxies with the corresponding, well-defined relations of the SMC (\citeauthor{STR:09}, but revised here in Sect.~\ref{sec:2}). Fundamental-mode (P0) as well as first-overtone (P1) Cepheids are considered. In addition, we consider the Cepheids in Sextans~A and Sextans~B (joined here into one set) and in Leo~A. The metallicity of the young population in these galaxies is still lower by a factor of three to four than in SMC. The question is whether this additional underabundance has a noticeable effect on the P-L and P-C relations. The most metal-poor galaxy known, i.e.\ I\,ZW\,18 with [O/H]$_{T_{\rm e}}=7.2$ \citep{Skillman:Kennicutt:93}, is not considered here because so far only a single Cepheid is known in the useful period range \citep{Fiorentino:etal:10}. The metallicities in the $T_{e}$-based system of \citet{Zaritsky:etal:94} of the galaxies in the present sample and their Galactic foreground reddenings \citep[from][]{Schlegel:etal:98} are given in Table~\ref{tab:01}. All data are corrected in the following for foreground reddening and absorption. \begin{table*} \centering \caption{Metallicities and foreground reddening of the sample galaxies.} \label{tab:01} \small \begin{tabular}{llll} \hline \hline \noalign{\smallskip} & \multicolumn{1}{c}{[O/H]$_{T_{\rm e}}$} & \multicolumn{1}{c}{Source} & \multicolumn{1}{c}{$E(B\!-\!V)_{\rm Gal}$} \\ \noalign{\smallskip} \hline \noalign{\smallskip} SMC & 7.98 & \citealt{Sakai:etal:04} & variable \\ IC\,1613 & 7.86 & \citealt{Sakai:etal:04} & 0.025 \\ WLM & 7.74 & \citealt{Sakai:etal:04} & 0.037 \\ Pegasus & 7.92 & \citealt{Skillman:etal:97} & 0.066 \\ Sex A \& B & 7.52 & \citealt{Skillman:etal:89} & 0.044, 0.032 \\ Leo A & 7.38 & \citealt{Skillman:etal:89,vanZee:etal:06} & 0.021 \\ \noalign{\smallskip} \hline \end{tabular} \end{table*} The P-C and P-L relations of the P0 and P1 Cepheids in the five sample galaxies and their distances are discussed in Sects.~\ref{sec:3}-\ref{sec:7}. The mean P0 and P1 distances are discussed in the light of independent distance determinations in Sect.~\ref{sec:8}. In Sect.~\ref{sec:9}, we compare the P-C and P-L relations of the metal-poor sample galaxies with the corresponding relations for more metal-rich Cepheids.

\label{sec:9} The P0 Cepheids in IC\,1613, WLM, Pegasus, Sextans~A \& B, and Leo~A (excluding SMC) are combined into composite P-L relations in $B$, $V$, and $I$, adopting the respective Cepheid distances derived in Sects.~\ref{sec:3}-\ref{sec:7} (Fig.~\ref{fig:09}). The resulting P-L relations, whose equations are at the foot of Fig.~\ref{fig:09}, are indistinguishable from those for SMC. Over the period interval of $0.2<\log P<1.2$, the P-L relations of SMC and the five sample galaxies agree to better than $0.02\mag$ in $V$ and $I$. In $B$, with fewer variables the agreement is not quite as good. In addition, the P1 Cepheids define closely agreeing P-L relations for SMC and the combined sample of five galaxies. This proves -- in agreement with our prediction -- that the P-L relations of SMC hold for the equally metal-poor galaxies IC\,1613, WLM, and Pegasus and even for the still more metal-poor Sextans~A \& B and probably also for Leo~A. (In the case of Leo~A, the comparison is restricted to Cepheids with $\log P<0.4$). The low-metallicity galaxies are therefore part of a family with (nearly) equal P-L relations. This holds of course also for the P-C relations, which are nothing else but the difference of the respective P-L relations. Cepheids of higher metallicity, such as those in LMC and the Galactic Cepheids in the solar neighborhood, have distinctly different P-L and P-C relations. For convenience, the coefficients of the relevant equations for the P0 Cepheids are compiled here in Table~\ref{tab:05} following the scheme $x=a\log P + b$. The equations for SMC and LMC follow from Sect.~\ref{sec:2}. The Galactic equations come from \citeauthor{TSR:03} and the revision in \citeauthor{STR:04}. \begin{table*} \begin{center} \caption{Coefficients of the relevant P-C and P-L relations for P0 Cepheids. Slope coefficients that agree to within $1\sigma$ are underlined.} \label{tab:05} \footnotesize \begin{tabular}{lrrrrcrrrrcrr} \hline \hline \noalign{\smallskip} & \multicolumn{4}{c}{SMC$^{1)}$} & & \multicolumn{4}{c}{LMC$^{2)}$} & & \multicolumn{2}{c}{Galaxy} \\ & \multicolumn{4}{c}{[O/H]=7.98} & & \multicolumn{4}{c}{[O/H]=8.36} & & \multicolumn{2}{c}{[O/H]=8.62} \\ & \multicolumn{2}{c}{$\log P<0.55$} & \multicolumn{2}{c}{$\log P>0.55$} & & \multicolumn{2}{c}{$\log P<0.9$} & \multicolumn{2}{c}{$\log P>0.9$} & & \multicolumn{2}{c}{} \\ & \multicolumn{1}{c}{$a$} & \multicolumn{1}{c}{$b$} & \multicolumn{1}{c}{$a$} & \multicolumn{1}{c}{$b$} & & \multicolumn{1}{c}{$a$} & \multicolumn{1}{c}{$b$} & \multicolumn{1}{c}{$a$} & \multicolumn{1}{c}{$b$} & & \multicolumn{1}{c}{$a$} & \multicolumn{1}{c}{$b$} \\ \noalign{\smallskip} \hline \noalign{\smallskip} $(B\!-\!V)^{0}$ & $0.191$ & $0.339$ & \underline{$0.415$} & $0.188$ && $0.306$ & $0.330$ & \underline{$0.435$} & $0.199$ && $0.366$ & $0.361$ \\[-2pt] & $\pm0.021$ & $\pm0.005$ & $\pm0.020$ & $\pm0.018$ && $\pm0.020$ & $\pm0.012$ & $\pm0.029$ & $\pm0.036$ && $\pm0.015$ & $\pm0.013$ \\[2pt] $(V\!-\!I)^{0}$ & $0.166$ & $0.511$ & \underline{$0.311$} & $0.413$ && $0.201$ & $0.474$ & \underline{$0.345$} & $0.331$ && $0.256$ & $0.497$ \\[-2pt] & $\pm0.018$ & $\pm0.005$ & $\pm0.016$ & $\pm0.014$ && $\pm0.017$ & $\pm0.010$ & $\pm0.024$ & $\pm0.030$ && $\pm0.017$ & $\pm0.016$ \\[2pt] $M_{B}^{0}$ & $-3.007$ & $-0.728$ & \underline{$-2.091$} & $-1.306$ && $-2.491$ & $-1.083$ & \underline{$-2.021$} & $-1.576$ && $-2.692$ & $-0.575$ \\[-2pt] & $\pm0.076$ & $\pm0.022$ & $\pm0.071$ & $\pm0.063$ && $\pm0.067$ & $\pm0.040$ & $\pm0.100$ & $\pm0.123$ && $\pm0.093$ & $\pm0.107$ \\[2pt] $M_{V}^{0}$ & $-3.203$ & $-1.071$ & \underline{$-2.531$} & $-1.466$ && $-2.787$ & $-1.414$ & \underline{$-2.505$} & $-1.713$ && $-3.087$ & $-0.914$ \\[-2pt] & $\pm0.060$ & $\pm0.018$ & $\pm0.056$ & $\pm0.050$ && $\pm0.048$ & $\pm0.029$ & $\pm0.074$ & $\pm0.091$ && $\pm0.085$ & $\pm0.098$ \\[2pt] $M_{I}^{0}$ & \underline{$-3.374$} & $-1.577$ & \underline{$-2.842$} & $-1.872$ && $-3.008$ & $-1.880$ & \underline{$-2.812$} & $-2.076$ && \underline{$-3.348$} & $-1.429$ \\[-2pt] & $\pm0.046$ & $\pm0.013$ & $\pm0.043$ & $\pm0.038$ && $\pm0.032$ & $\pm0.019$ & $\pm0.057$ & $\pm0.069$ && $\pm0.083$ & $\pm0.097$ \\ \noalign{\smallskip} \hline \end{tabular} \end{center} \tablefoot{% $^{1)}$ adopted at $(m-M)^{0}_{\rm SMC}=18.93$ (\citeauthor*{TSR:08b}, Table~6); $^{2)}$ adopted at $(m-M)^{0}_{\rm LMC}=18.52$ (\citeauthor*{TSR:08b}, Table~7) } \end{table*} The steep slopes of the Galactic P-L relations from Paper~I and II corresponds to data for Cepheids in Galactic clusters and OB associations \citep{Feast:99} as well as Baade-Wesselink-Becker distances by \citet{Fouque:etal:03} and \citet{Barnes:etal:03}, the two fully independent methods leading to the same result. Criticism of the result was discussed by \citeauthor*{TSR:08b}. The P-L relations of metal-rich Cepheids will be discussed in more detail in a forthcoming paper; it is possible that they experience a break at long periods ($\log P\ga1.6$), but this is irrelevant here. The observed P-L relations of LMC are closely matched -- including the break at $\sim\!10^{\rm d}$ -- by theoretical P-L relations based on pulsation models \citep{Marconi:etal:05}. The same models do not show a break at higher metallicities in agreement with the Galactic P-L relations adopted here. \citet{Marconi:etal:10} also derived theoretical P-L relations for ultra-low metallicities. They have no break and are somewhat flatter than in SMC up to $\log P=0.55$, but are much steeper beyond that point. The comparison may not be justified because the adopted metallicity ([O/H]$\sim 7.0$) is lower than in SMC and even Leo~A. To illustrate the difference between the Cepheids in SMC, LMC, and the solar neighborhood, the P-C and P-L relations of SMC and LMC are plotted {\em relative\/} to those of the Galaxy in Fig.~\ref{fig:10}. In each panel, the Galactic relations are taken as reference and the {\em differences\/} in color and absolute magnitude of the Cepheids in the other two galaxies are shown as a function of $\log P$ (in the sense $x_{\rm LMC/SMC}-x_{\rm Galaxy}$). \begin{figure*} \centering \resizebox{0.7\hsize}{!}{\includegraphics{fig10.eps}} \caption{a) The P-C relation in $(B\!-\!V)$ of P0 Cepheids in both the very metal-poor SMC and the relatively metal-poor LMC {\em relative\/} to the metal-rich Solar neighborhood. b) Same as a) but for $(V\!-\!I)$. c) - e) The P-L relations in $B$, $V$, and $I$, respectively, of the P0 Cepheids in SMC and LMC {\em relative\/} to the Solar neighborhood. The artificial spikes of the relations are due to statistical errors of the fits below and above the break. The lines are only drawn over the period range where they are well defined by observations.} \label{fig:10} \end{figure*} As seen in Fig.~\ref{fig:10}a the LMC Cepheids are bluer in $(B\!-\!V)$ than their Galactic counterparts by up to $0.09\mag$ at $\log P=0.9$. The color difference is even larger between SMC and the Galaxy, i.e.\ $0.13\mag$ at $\log P=0.55$. The red color of the Galactic Cepheids is due to their lower temperature and the blanketing effect of the metal lines (see \citeauthor{STR:04}). The color differences in $(V\!-\!I)$ between the three galaxies in Fig.~\ref{fig:10}b are smaller. The LMC Cepheids are bluer than in the Galaxy by up to $0.08\mag$. Unexpectedly, the SMC Cepheids are redder than in LMC, yet still bluer than those in the Galaxy by a marginal amount of $0.04\mag$ or less, depending on period. The P-L relations in $B$, $V$, and $I$ of LMC and SMC are plotted relative to the Galactic P-L relations in Fig.~\ref{fig:10}c-e. The relations of LMC and SMC have similar characteristics and differ mainly in the zero-point, but they are both much flatter than in the Galaxy beyond the break point. At $\log P=0.6$, LMC and SMC Cepheids are respectively brighter by $0.39\mag$ and $0.37\mag$ than in the Galaxy, whereas at $\log P = 1.7$ they are fainter by 0.14 and $0.29\mag$, respectively. LMC Cepheids are brighter than in SMC by $0.15-0.20\mag$, somewhat depending on period. The significant luminosity differences of the Cepheids in the three galaxies cannot be explained by errors in the adopted distances, which are on the order of $0.1\mag$. In addition, it is impossible to explain the different slopes of the P-L relations by distance errors. We note that some of the {\em slopes\/} of the P-L relations in Table~\ref{tab:05} show striking agreement. The slopes of SMC and LMC are essentially identical in $B$, $V$, and $I$ above the break points, and the short-period SMC P-L relation in $I$ has the same slope as Galactic Cepheids. In addition, the slopes of the P-C relations of SMC and LMC are the same to within $\sim\!1\sigma$ for $\log P>0.9$. The Cepheids designated here as low-metallicity objects comprise in fact a wide metallicity range of $8.0>$[O/H]$T_{\rm e}>7.4$. Their very similar P-L relations imply that they are quite insensitive at these low levels to metallicity changes. This is in sharp contrast to more metal-rich Cepheids where a change of only $\Delta$[O/H]$T_{\rm e}=0.26$ causes the pronounced differences between the LMC and Galactic P-L relations. The use of Cepheids as distance indicators has been extended here to include fundamental-mode (P0) and first-overtone (P1) Cepheids with the shortest periods known. Among the known Cepheid population of the SMC, 47\% of the P0 pulsators have periods less than $\log P=0.4$, extending down to $\log P=0.0$, and 37\% are P1 pulsators with periods down to $\log P=-0.2$. The large number of these additional Cepheids makes them indispensable for accurate distance determinations. The distances derived here agree with independent RR~Lyr and TRGB distances to within a few $0.01\mag$.

1012.4940

1012

1012.3084_arXiv.txt

We outline the key--steps towards the construction of a physical, fully relativistic cosmology, in which we aim to trace Dark Energy and Dark Matter back to physical properties of space. The influence of inhomogeneities on the effective evolution history of the Universe is encoded in backreaction terms and expressed through spatially averaged geometrical invariants. These are absent and interpreted as missing dark fundamental sources in the standard model. In the inhomogeneous case they can be interpreted as energies of an emerging scalar field (the morphon). These averaged invariants vanish for a homogeneous geometry, where the morphon is in an unstable equilibrium state. If this state is perturbed, the morphon can act as a classical inflaton in the Early Universe and its de--balanced energies can mimic the dark sources in the Late Universe, depending on spatial scale as Dark Energy or as Dark Matter, respectively. We lay down a line of arguments that is qualitatively conclusive, and we outline open problems of quantitative nature, related to the interpretation of observations.

\subsection{The foliation issue and the notion of an effective cosmology} The homogeneous--isotropic standard model of cosmology, being itself a particular solution of Einstein's general theory of relativity, does by far not exploit the degrees of freedom inherent in the geometry as a dynamical variable. It is this richer tone of general relativity -- as compared to the Newtonian theory -- that opens the possibility to generalize cosmological models, notably by including inhomogeneous structure also in the geometrical variables. There are several guidelines to be emphasized in such a generalization: firstly, a cosmology is thought of as an evolving space section that implies the need to speak of a foliated space time, introducing four degrees of freedom (the lapse and shift functions in an ADM setting). This necessarily implies, on general grounds, a breaking of four--dimensional covariance. This fact should not be confused with coordinate-- or gauge--dependence of the resulting cosmological equations and variables, however. Secondly, a cosmology purports an effective point of view in the sense that the evolving spatially inhomogeneous variables are thought of as being ``averaged over'' in a way that has to be specified. We aim at a description that only implicitly refers to a metric. However, if a metric is to be specified, a cosmological metric is then to be considered as an effective, ``smoothed out'' or {\it template metric}, being not necessarily a solution of the equations of general relativity. Finally, a {\it physical} cosmology should be characterized by such an effective evolution model, an effective metric to provide the distance scale for the interpretation of observations, or alternatively an evolution model for average characteristics on the light cone, together with a set of initial data. These latter are to be related to physical properties of fundamental sources, but also to the geometrical data at some initial time (effective, i.e ``averaged'' quantities of known energy sources, intrinsic and extrinsic curvature). This latter clearly emphasizes the absence of any phenomenological parameters. Those would just parametrize our physical ignorance. All these points will be made explicit in what follows. \subsection{The dark side of the standard model: postulated sources and proposed solutions} The high level of idealization of the geometrical properties of space in the standard model leads to the need of postulating sources that would generate ``on average'' a strictly, i.e. globally and locally, homogeneous geometry. It is here where a considerable price has to be paid for a model geometry that obviously is not enough to meet physical reality, unless we really believe that we can find the missing sources: $96$ percent of the energy content is missing in the form of a) a postulated source acting attractive like matter, so--called Dark Matter ($\cong 23$ percent) and b) a postulated source acting repulsive, so--called Dark Energy ($\cong 73$ percent). Evidence for the former does indeed come from various scales (galaxy halos, clusters and cosmological, see e.g. \cite{roos}), while evidence for the latter only comes from the apparent magnitude of distant supernovae (see \cite{SNIa:Union,SNIa:Constitution,SNIa:Essence} for the latest data) that, if interpreted within standard model distances, would need an accelerating model. In the simplest case this volume acceration is achieved by a homogeneous--isotropic cosmology with a cosmological constant. It should be emphasized that when we speak of evidence, we already approach this evidence with model priors \cite{huntsarkar,seikel:acc,cmbobs}. Keeping this idealization for the geometry of the cosmological model for example, one has to conjecture fundamental fields in proportion to the missing dark components on cosmological scales. The search for these fields is one major research direction in modern cosmology. \smallskip Another huge effort is directed towards a generalization of the underlying theory of gravitation. While this would generalize the geometry of the model, it is not clear why all these efforts go into a generalization of general relativity and not into the generalization of the cosmological model within general relativity. There are certainly good lines of arguments and various motivations in particle physics and quantum gravity to go beyond the theory of Einstein (for reviews see \cite{DE:review}, \cite{DE:pilar}), but the ``dark problem'' may be first a classical one. \smallskip Looking at generalizations of the standard model \emph{within} general relativity can be identified as a third research direction to which we dedicate our attention here. In light of current efforts it is to be considered conservative, since it does not postulate new fundamental fields and it does not abandon a well--tested theory of gravitation \cite{dressing}, \cite{rasanen:de}, \cite{kolb:backreaction} (for reviews on the physical basis of this third approach see \cite{buchert:jgrg,buchert:review} and \cite{rasanen:acceleration}). Among the works in this latter field, research that analyzes spherically symmetric exact solutions has been meanwhile developed to some depth, and has determined the constraints, necessary to explain Dark Energy, on a postulated observer's position within a large--scale void (see \cite{LTB:review,bolejkoandersson,celerier,voidtest} and references therein). \subsection{Fictitious and physical backgrounds} Perhaps a reason for not questioning the standard model geometry within general relativity and to go for the search for fundamental fields or for generalizations of the laws of gravitation is the following belief: effectively, i.e. ``on average'', the model geometry has to be {\it homogeneous}, since structures should be ``averaged over''. Then, due to observational facts on large scales (the high degree of isotropy of the Cosmic Microwave Background, if the dipole is completely eliminated due to our proper motion with respect to an idealized exactly isotropic light sphere), and first principle priors (the {\it strong cosmological principle} that requires the universe model to look the same in all directions), the model geometry is taken to be {\it locally isotropic}. \smallskip Taking this reasoning at face value we must note two points: the notions of homogeneity and isotropy in the standard model are both too strong to be realistic: firstly, local isotropy implies a model that is locally and globally homogeneous, i.e. despite the conjecture that the homogeneous model describes the inhomogeneous Universe ``on average'', this {\it strict homogeneity} does not account for the fact that any averaging procedure, in one way or another, would introduce a {\it scale--dependence} of the averaged (homogeneous) variables \cite{ellisbuchert}. This scale--dependence, inherent in any physical averages, is suppressed. Even if a large {\it scale of homogeneity} exists (we may call this {\it weak homogeneity principle}), the model is in general scale--dependent on scales below this homogeneity scale \cite{sylos:copernican}. The same is true for isotropy: while the averaged model may be highly isotropic on large scales, a realistic distribution on smaller scales is certainly not (we may call this {\it weak isotropy principle}). Correspondingly, a {\it weak cosmological principle} would be enough to cover the reality needs while still facing observational evidence on large scales. \smallskip We may summarize the above thoughts by noting that, on large scales, a homogeneous--(almost)isotropic {\it state} does not necessarily correspond to a homogeneous--(almost)isotropic {\it solution} of Einstein's equations. These former states are the averages over fluctuating fields and it is only to be expected that the state coincides with a strictly homogeneous solution in the case of absence of fluctuations. In other words, looking at fluctuations first requires to establish the average distribution. Only then the notion of a {\it background} makes physical sense \cite{kolb:backgrounds}. Current cosmological structure formation models, perturbation theories or N--body simulations, are constructed such that the average vanishes on the background of a homogeneous--isotropic {\it solution} \cite{buchertehlers}. A such chosen reference background may be a {\it fictitious background}, since it arises by construction rather than derivation. On the contrary, a {\it physical background} is one that corresponds to the average (whose technical implementation has to be specified, and which is nontrivial if tensorial quantities like the geometry have to be ``averaged''). A sound implementation of a physical background will be a statistical background where not only solutions but ensembles of solutions are averaged. Having specified such an averaging procedure, a physical cosmological model may then be defined as an evolution model for the average distribution. Despite these remarks it is of course possible that the homogeneous solution forms at the same time the average. A well--known example is Newtonian cosmology \cite{buchertehlers}. It is also conceivable that the homogeneous solution provides, in some spatial and temporal regimes, a good approximation for the average. Still, it is important to consider perturbations on the correct background solution \cite{kolb:voids}.

1012.3084

1012

1012.2800_arXiv.txt

We present the N, O, F and Na abundance and $^{12}$C/$^{13}$C isotopic ratio measurements or upper limits for a sample of 10 C-rich, metal-poor giant stars, eight enhanced in s-process (CEMP-s) elements and two poor in n-capture elements (CEMP-no). The abundances are derived from IR, K-band, high-resolution CRIRES@VLT spectra obtained. The metallicity of our sample ranges from [Fe/H]=$-$3.4 to $-$1.3. F abundance could be measured only in two CEMP-s stars. With [F/Fe]$=$0.64, one is mildly F-overabundant, while the other is F-rich, at [F/Fe]$=$1.44. For the remaining eight objects, including both CEMP-no in our sample, only upper limits on F abundance could be placed. Our measurements and upper limits show that there is a spread in [F/C+N] ratio in CEMP-s stars as predicted by theory. Predictions from nucleosynthetic models for low-mass, low-metallicity Asymptotic Giant Branch stars, account for the derived F abundances, while the upper limits on F content derived for most of the stars are lower than the predicted values. The measured Na content is accounted for by AGB models in the 1.25 to 1.75\,M$_{\odot}$ range, confirming that the stars responsible for the peculiar abundance pattern observed in CEMP-s stars are low-mass, low-metallicity AGB stars, in agreement with the most accepted astrophysical scenario. We conclude that the mechanism of F production in current state-of-the-art low-metallicity low-mass AGB models needs further scrutiny and that F measurements in a larger number of metal-poor stars are needed to better constraint the models.

The understanding of the stellar nuclear production sites and evolution scenarios is of great current interest. A large number of high-resolution spectral studies targeting metal-poor objects selected by mining the HK survey \citep{beers92} and the HES \citep{christlieb01} have measured a wide set of chemical elements ranging from Li \citep[e.g.][]{spite05} to the n-capture species \citep[e.g.][]{sneden03}, providing constraints on how the stars formed as well as on the nucleosynthetic processes that took place in the early Galaxy. In contrast, only one single measurement of F in extremely metal-poor (EMP) stars has been published to date \citep{schuler07}. The most accessible lines suited for F measurements, the vibration-rotation transition of the HF molecule, are located in the K band, requiring the use of high-resolution, IR spectrographs mounted on 8\,m class telescope to target halo stars, a configuration not available until recently. Fluorine is an element of particular interest, extremely sensitive to the physical conditions within stars. Although the single stable F isotope, $^{19}$F, is not involved in the main reactions taking place in the cores of stars, it can be created and destroyed in several different ways, and its dominant source is not yet clear. Theoretical modeling has indicated (at least) three possible sites for F production. \citet{woosley98} proposed first that F is produced by neutrino spallation on $^{20}$Ne; $^{20}$Ne($\nu$,$\nu ' p$)$^{19}$F. Second, modeling shows that low- metallicity, low-mass (M$\leq$3-4\,M$_{\odot}$) AGB stars synthesize very large quantities of F \citep[see e.g.][]{cristallo07} via the reactions $^{14}$N(n,p)$^{14}$C($\alpha$,$\gamma$)$^{18}$O($p,\alpha$)$^{15}$N($\alpha,\gamma$)$^{19}$F {\bf and $^{14}$N($\alpha$,$\gamma$)$^{18}$F($\beta^{+}$)$^{18}$O(p,$\alpha$)$^{15}$N($\alpha$,$\gamma$)$^{19}$F}, where the neutrons are provided by $^{13}$C($\alpha,n$)$^{16}$O and the protons mainly by $^{14}$N($n,p$)$^{14}$C \citep[see][]{jorissen92}. On the other hand, massive AGB stars are expected to destroy F via Hot Bottom Burning (HBB) \citep[see e.g.][]{smith05,karakas07}. Finally, \citep{meynet00} proposed that Wolf-Rayet stars are a possible source for F, through the $^{14}$N($\alpha,\gamma$)$^{18}$ F($\beta^{+}$)$^{18}$O($p,\alpha$)$^{15}$N($\alpha,\gamma$)$^{19}$F chain. While the physical conditions for F production exist in several phases of stellar evolution, the precise contribution of each of these three sites to the Galactic F evolution is not known, The measurements of F content in any object have been few and give conflicting evidence about the origin of F in the Universe. \citet{renda04} showed that the inclusion of all of sites in Galactic chemical evolution models is necessary to reproduce the F abundances measured in Milky Way field stars. In their model, AGB stars are responsible for significant amounts of F production in the early Universe, because of the metallicity dependence of their yields, while WR stars are significant contributors at solar and super-solar metallicities. \citet{jorissen92} derived F abundances in giants that recorded the chemical evolution (K and M giants), in AGB giants that were dredging up freshly minted F from their interiors (MS, S, SC, N and J giants), and in giants that had been polluted by one of the former (Ba giants). They suggested that the overproduction of F in the AGB stars (factors of 3 to 30 with respect to the solar system in the most extreme S and N stars) showed that AGB stars are likely to major contributors to the Galactic F abundance, at least at metallicities close to solar. In addition, the largest F overabundances measured could not be explained with standard AGB models and required additional mixing to achieve the desired amounts. However, recent results by \citet{abia08} and \citet{abia10} suggest that because of a possible lack of proper accounting for C-bearing molecules (i.e. CH, CN, CO and C$_2$) contribution, the F abundances reported in \citet{jorissen92} for solar metallicity giants, which relied on the same HF line used in the present analysis, had been overestimated. Therefore, the long standing problem of the discrepancy between the high [F/Fe]\footnote{ Hereafter:$\log \epsilon$(A) = log N(A)+12. [X/Y] = log(X/X$_{\odot}$) $-$ log(Y/Y$_{\odot}$).} and [F/O] ratios measured in solar metallicity giants and the low ones resulting from the models, is likely solved by the adoption of a more complete molecular line list for the synthetic spectra. In fact, the \citet{cristallo07} models account fully for the F abundances measured in several C-stars ranging from solar metallicity to slightly metal-poor ([Fe/H]$\sim$0.0 to $\sim$-0.5) \citep{abia08,abia10}. While the above results make it clear that solar metallicity AGB stars make at least some F, there is not yet a consensus on the overall contribution of AGB stars throughout the chemical evolution of stellar systems. \citet{cunha03} measured F in red giants in the LMC and $\omega$ Cen. They found that F/O declines as the O abundance decreases and the two $\omega$ Cen giants have particularly low F/O values. They argue that their results are qualitatively consistent with most F production coming from either neutrino nucleosynthesis or WR stars rather than AGB stars in these systems, which certainly have a different chemical evolution than the Milky Way. In particular, the low F abundances in metal-poor ([Fe/H]$\sim -1$) $\omega$ Cen stars that are enriched in s-process elements is difficult to reconcile with the idea that metal-poor AGB stars are major contributors to F production. F has also been measured in stars in the Galactic bulge \citep{cunha08}. The F/O results there can be explained by contributions by both AGB and WR stars, although the lack of an s-process enhancement in the most F-rich bulge stars suggests that AGB stars may play a less prominent role in the bulge than has been inferred for the disk. Measurements of fluorine in the interstellar medium \citep{federman05} show no evidence of F over-abundances due to the neutrino process in SNII. The F production in rotating WR stars has been reconsidered by \citet{palacios05}, who found that F yields are significantly lower than the \citet{meynet00} predictions, indicating that their contribution to the Galactic F budget would be negligible. These results suggest that at low-metallicity AGB stars play a major role in F production, although this idea would leave the $\omega$ Cen results unexplained. The idea of AGB stars as producers of F is supported by the large F enhancements found in post-AGB stars \citep{werner05} and planetary nebulae \citep[see e.g.][and references therein]{otsuka08}, the progeny of AGB stars as well as the results for abundances in Milky Way AGB stars cited above. However, these observations are focused on [Fe/H]$\sim$0 metallicity stars and do not address directly the question of the production of F production in metal-poor AGB stars. While Na measurements have been obtained for a much larger number of metal-poor stars, no quantitative study has been performed so far on the role of AGB stars in its production, in particular at low metallicity. In thermally pulsing AGB stars, Na is mostly synthesized by proton captures on $^{22}$Ne. Models indicate that the main source of Na in low-mass AGB stars, at least at solar metallicity, is the creation of a Na pocket, located at the top (or near the top) of the $^{14}$N pocket. At lower metallicities, other mechanisms, such as the neutron capture on $^{22}$Ne, which can occur both during the radiative $^{13}$C burning and during the convective $^{22}$Ne burning, may become important (Cristallo et al 2006b and references therein). In general, according to e.g., Cristallo et al. and \citet{bisterzo10}, the production of Na increase by about 1\,dex or more going from solar to [Fe/H] ~ -2.3, a similar metallicity to that of the stars in our sample. This is due to a higher efficiency of all the Na nucleosynthesis channels described above. The testing of such predictions is important to constraint details of the key reactions involved. Carbon-enhanced metal-poor (CEMP) stars provide an opportunity to directly measure the F and Na production in low-mass, metal-poor AGB stars. These stars are chemically peculiar objects, characterized by an overabundance of C ([C/Fe]$>$1\footnote{ Some authors use different cutoffs values depending on the evolutionary state of the star, with cutoff enhancement for giants as low as [C/Fe]$\simeq$0.5 \citep[see][]{aoki07}}) accounting for 10-20\% of stars below [Fe/H]$\leq-$2.5 \citep{marsteller05,cohen05,lucatello06}. Less than a third of CEMP stars exhibit no enhancement in heavy elements (CEMP-no), while most of these objects \citep[over 70\% ][]{aoki07} are characterized, by an overabundance of n-capture, s-process, elements (CEMP-s). \citet{lucatello05} showed that likely {\it all} C-rich, extremely metal-poor stars with s-process enhancement (the CEMP-s stars) belong to a binary system. CEMP-s are then the metal-poor analog to the classical CH and Ba stars: low-mass stars (M$\sim$0.8\,M$_{\odot}$) whose slightly more massive (between $\simeq$1.2 and $\simeq$2.5\,M$_{\odot}$, the exact range depending on metallicity) companion, now a faint white dwarf, dumped material processed during its AGB phase on their surfaces, leaving its chemical fingerprints in the composition of their envelopes. Therefore, the nucleosynthetic processes taking place in extremely metal-poor, low-mass ($\simeq$1.5\,M$_{\odot}$) stars, now long extinct, can be investigated through the study of CEMP-s stars characteristics \footnote{Because of its high Eu abundance, not fully accountable with standard s-process nucleosynthesis, one of the stars in our sample, HD187861, has been included in the CEMP-rs category. However, since all the likely formation scenarios invoked to explain the abundance patterns of these objects include mass transfer from an AGB companion \citep[see e.g.][]{jonsell06,lucatello09}, as far as the present discussion is concerned HD187861 can be considered as an CEMP-s}. The origin of CEMP-no objects is, on the other hand, still a mystery. \citet{fujimoto00} suggested that they may have formed as chemically ``normal'' low-mass stars and became C-enhanced through a path of self-enrichment due to anomalous mixing processes specific to low-metallicity stars. Alternatively, as \citet{ryan05} suggested, they could have been born from C-rich gas, possibly polluted by a previous generation of supernovae whose fall-back avoids the ejection of heavier elements during the explosion \citep[e.g.,from high-mass, rapidly rotating stars see][and references therein]{meynet06}. Objects like the recently discovered extremely-metal poor, C-rich Damped Lyman-$\alpha$ system \citep{cooke10} might turn out to be the connection between the yields of the Pop III stars and their later incorporation into CEMP-no stars. On the other hand, their abundance patterns could arise from early AGB transfers from low-mass stars before any considerable s-process element production took place \citep{ryan05,masseron09a}, or alternatively from an AGB stars whose evolution was truncated by binary interaction (see, e.g., Izzard \& Tout 2003) with its companion, the presently observed object. The measurement of a value or a strong upper limit on F abundance is of crucial importance to probe the origin of the chemical pattern observed in CEMP-no stars given that F can be synthesized before the bulk of the s-process-element production. In \citet{cristallo07}, for instance, as early as the third-dredge-up episode for a 2\,M$_{\odot}$ at [Fe/H]=$-$2.3, [F/C]=-0.4 and [F/Fe]=1.65, [Ba/C]=-1.5 and [Ba/Fe]=0.5 (and [C/Fe]=2.05). Therefore, a measurement in CEMP-no stars of an amount of F comparable to C would strongly argue in favor of AGB enrichment. \citet{schuler07} measured the F abundance in a extremely metal-poor star for the first time, deriving an abundance of $\log \epsilon$(F)= $+4.96 \pm 0.21$ corresponding to an abundance ratio [F/Fe] = $+$2.9 for the CEMP-s star HE~1305+0132. \citet{lugaro08} compared this value to existing nucleosynthesis and mass transfer models. Conclusion that an object with such an extreme F content should be exceedingly rare, while most CEMP-s stars are expected to exhibit a noticeable but smaller F overabundance. We here present F abundance measurements for two CEMP-s and upper limits for eight more CEMP stars (six CEMP-s and two CEMP-no) and discuss the implications of our result on our current understanding of AGB nucleosynthesis.

We have obtained IR high-resolution observations in a sample of ten CEMP stars, eight CEMP-s and two CEMP-no. The aim was to measure Na abundances, $^{19}$F contents from HF, O from CO and N from CN lines in the $K$ band. A comparison of predictions from different families of low-metallicity, low-mass (1.25 to 1.75 M$_{\odot}$ range) AGB nucleosynthesis models show that they reproduce well the observed Na abundances in the CEMP-s stars in our sample. This result strongly argues in favor of a polluter of low-mass for these objects, in agreement with the currently most accepted scenario for their formation. F could only be measured in two CEMP-s stars, while for the remaining eight objects only upper limits could be derived. However, these data are sufficient to show that a range of [F/C+N] values are produced in low-metallicity AGB stars, in accord with predictions of the mass- and metallicity-dependence of F production in AGB models. The only two F measurements obtained (the CEMP-s stars HE\,1152-0355 and HD\,5223 respectively at [Fe/H]=$-$1.27 and $-$2.06) are accounted for, within the errors, low-mass, low-metallicity AGB models. On the other hand, most of the derived upper limits for F abundance in CEMP-s are not satisfactorily accounted for by nucleosynthetic computations. In fact, the comparison with four of the most recent models for low-mass (2.0\,M$_{\odot}$), low-metallicity AGB nucleosynthetic models shows that there are large differences in the predictions between different families of models, which cannot reproduce several of the upper limits, not providing any [F/C+N] ratios predictions low enough to account for the values measured in several of the sample stars. A comparison with \citet{karakas07} models for different masses indicates that only objects more massive (2.5\,M$_{\odot}<$M$<$4\,M$_{\odot}$) that those generally considered as responsible for CEMP-s peculiar abundances (M$\simeq$1.5\,M$_{\odot}$) produce the [F/C+N] ratios observed in about two thirds of our CEMP-s sample. This possibility is not only unlikely because of simple IMF considerations, but it is also challenged by the fact that at low metallicity, HBB is active in stars in this mass range, producing N and depleting C (bringing the ratio to [N/C]$\sim$1 \citep[see e.g.][]{johnson07}) and preventing the star from becoming a CEMP-s polluter. {\bf As shown in Fig \ref{fig4}, the \citet{karakas07} 1.25\,M$_{\odot}$ model is the low-mass AGB ones that comes closest to explaining the upper limits. Given the discussed spread present in F production in different models, it is indeed possible that it affects also predictions for other masses, including 1.25\,M$_{\odot}$ , hence accounting for the derived upper limits. It is however impossible, at present time, to systematically test this differences, as only \citet{karakas07} provide F production calculations for several AGB masses.} We discussed two possible solutions that could explain the lowest F upper limit range. One is that the standard evolution of an AGB star may be truncated by binary interaction, and a lower F abundance can be obtained in early TP pulses (even though this would lead to lower Ba production, which could make the accounting for observed Ba abundances possibly more challenging). The other is the action of CBP, which may reduce the F in the AGB envelope if such extra-mixing processes exposes material at temperatures high enough to activate $^{19}$F(p,$\alpha$). The conclusion that can be drawn from the comparison between the data and the current models is that F production in low metallicity AGB stars is probably not as high as expected on the basis of the current models. It is important to keep in mind that, as discussed in the introduction, at solar metallicity model predictions are in agreement with the observations, therefore such problem seem to be peculiar to the low metallicity regime. AGB nucleosynthesis, especially at low metallicity, is still not fully understood and a large number of uncertainties affect in particular F nucleosynthesis. Improvements in nuclear reaction rate (such as $^{14}$C($\alpha$,$\gamma$)$^{18}$O and $^{18}$F($\alpha$,p)$^{21}$Ne) accuracy are needed, and a better grasp of the CBP mechanism is highly desirable. Only one of the available models, by \citet{cristallo07} takes into account the appropriate C, N and O enhancements in the computation of the stellar opacities; also, the prescription for mass loss is still rather uncertain. Moreover, currently no model takes into account the effect on the AGB evolution of the presence of a (close) companion, which could stimulate extra mixing and/or truncate the AGB phase (see e.g. Izzard \& Tout 2003), affecting the nucleosynthesis. The inclusion of these ingredients into a new generation of AGB models might in the future be able to account for the present observations. From an observational point of view, pursuing a larger number of F abundance measurements in CEMP stars is extremely important. In particular, K-band high-resolution spectral observations of cool (T$_{eff}<$4200\,K) CEMP stars would allow actual measurements (rather than upper limit placements) of F even for low F abundances ($\log \epsilon$(F)$\sim$2). These measurements would provide more stringent constraints that are urgently needed for a future generation of AGB models and to effectively probe F chemical evolution.

1012.2800

1012

1012.0805_arXiv.txt

{The distances that galactic cosmic ray electrons and positrons can travel are severely limited by energy losses to at most a few kiloparsec, thereby rendering the local spectrum very sensitive to the exact distribution of sources in our galactic neighbourhood. However, due to our ignorance of the exact source distribution, we can only predict the spectrum \emph{stochastically}. We argue that even in the case of a large number of sources the central limit theorem is not applicable, but that the standard deviation for the flux from a random source is divergent due to a long power law tail of the probability density. Instead, we compute the expectation value and characterise the scatter around it by quantiles of the probability density using a generalised central limit theorem in a fully analytical way. The uncertainty band is asymmetric about the expectation value and can become quite large for TeV energies. In particular, the predicted local spectrum is marginally consistent with the measurements by Fermi-LAT and HESS even without imposing spectral breaks or cut-offs at source. We conclude that this uncertainty has to be properly accounted for when predicting electron fluxes above a few hundred GeV from astrophysical sources.}

The interest in the propagation of galactic cosmic ray (GCR) electrons and positrons from astrophysical sources, like supernova remnants (SNRs), has recently been revived, mostly in light of contemporary, partly ``anomalous'' measurements~\cite{Torii:2008xu,Adriani:2008zr,Chang:2008zzr,Collaboration:2008aaa,Abdo:2009zk,Aharonian:2009ah}. In the context of their possible explanation as exotic signals from dark matter annihilation or decay, it is usually believed that astrophysical sources of GCRs only lead to local power law spectra, in contrast to strong features from exotic sources. This however relies on a simplified picture of propagation assuming not only a steady state situation but also a continuous spatial distribution of sources. On the time and distance scales of GCR propagation, however, astrophysical sources can be considered \emph{discrete}. Diffusion renders the fluxes originating from point-like and from extended sources like old SNRs indistinguishable. Furthermore, although the details of the injection of particles from SNRs is not fully understood (for a discussion, see ref.~\cite{Caprioli:2009fv}), the bulk of the accelerated GCRs is expected to be released with the onset of the radiative phase, that is within a relatively short time. For nuclear GCRs, the spectrum depends rather weakly on the exact distribution of sources in space and time since protons and nuclei diffuse over distances of kiloparsecs before escaping from the cosmic ray halo, thereby averaging over the distribution of sources on these scales. It is therefore admissible to neglect the small scale distribution and to assume a steady and continuous distribution of sources. Such a distribution function can be obtained by generalising information from radio or x-ray surveys of source candidates~\mbox{\cite{Case:1998qg,Lorimer:2003qc}} and is being used in many computations of GCR fluxes including the most well-known numerical \cite{Moskalenko:1997gh,galprop,Evoli:2008dv} and semi-analytical codes \cite{Maurin:2001sj,Maurin:2002hw}. The propagation of leptonic GCRs is, however, dramatically different as electrons and positrons suffer from strong energy losses due to interactions with the galactic magnetic fields and interstellar radiation fields (ISRFs). Therefore, the diffusion-loss length, i.e. the distance electrons and positrons can travel away from the sources without loosing virtually all energy, is much shorter than for protons and nuclei. In addition, this distance is energy dependent such that at high energies, that is above $100 \, \text{GeV}$ or so, mostly young and nearby sources contribute and the spectrum at these energies strongly depends on the history and spatial distribution of sources within a kiloparsec. Therefore, it is \emph{not} admissible to neglect the nearby small scale distribution. It has been proposed~\cite{Kobayashi:2003kp} to model the electron-positron flux by considering a continuous distribution of sources for distances $\gtrsim 1 \, \text{kpc}$ or ages $\gtrsim 10^5 \, \text{yr}$, and to supplement it by the few known nearby SNRs. However, this requires that \emph{all} nearby sources are known from surveys which seems, at least, challenging. First of all, such studies suffer from various selection effects. Radio surveys, for instance, are insensitive to SNRs of low surface brightness and also young but distant ones~\cite{Green:2005yt}. Furthermore, for individual objects it can be difficult to decide whether they can contribute to the local electron-positron flux. For example, it has been suggested that pulsars can accelerate electrons and positrons but it is not clear if theses high energy particles can escape efficiently from the surrounding pulsar wind nebulae (PWNe). Recent observations~\cite{Bamba:2010zk}, however, suggest that this may be the case. In addition, the prediction of a flux requires accurate distance information. Considering the uncertainty in the distance estimates of some of the known SNRs (for example, for RX J1713.7-3946 distance estimates vary between $1$ and $6 \, \text{kpc}$ \cite{Koyama:1997wp,Slane:1999xr}) we cannot expect the distances of sources yet to be discovered to be much more reliable. Most importantly however, surveys using electromagnetic radiation will only tell us about sources on the past light cone. The propagation of charged GCRs is however dramatically different from the rectilinear propagation of photons such that much older sources can potentially contribute. Consequently, relying on surveys we might simply miss old but nearby sources which could lead to artificial features in the predicted electron-positron flux (for an illustration of this, see ref.~\cite{Ahlers:2009ae}). It thus seems difficult to determine the complete \emph{real} distribution of sources. However, as we do not expect temporal variations in the source rate to be too strong, we can generalise the information about sources on the past light cone to earlier times thereby building a model of the statistical distribution in both, position and time. The predictions possible using this statistical information are however only \emph{stochastic}, that is one can predict the expectation value for the flux from a statistical ensemble of sources. For protons and nuclei, the fluctuations around the expectation value are rather small but for electrons and positrons, the flux from a particular system in the statistical ensemble, that is a particular distribution of sources like the one in the Galaxy, will in general differ from the expectation value. In this sense our ignorance of the true distribution of sources in age and distance introduces an uncertainty into the predicted electron-positron flux, in particular at high energies. This uncertainty can in principle be investigated by Monte Carlo methods~\cite{Pohl:1998ug,Strong:2001qp,Swordy:2003ds}, that is by computing the fluxes from a large number of systems, randomly drawing the source positions and ages from the statistical model. Here we aim at a \emph{fully analytical} calculation which will allow us to check against possible prejudices, e.g. the gaussianity of the probability density for fluxes. Furthermore, this approach provides the transparency to trace the propagation of model parameters into the final result. The paper is organised as follows. In section~\ref{sec:Propagation} we briefly review the usual setup for propagation of GCR electrons and positrons in a purely diffusive model. The calculation for the expectation value of the flux is outlined in section~\ref{sec:Expectation} and the determination of the uncertainties is explained in section~\ref{sec:Uncertainties}. We present and discuss our results in section~\ref{sec:Discussion} and conclude in section~\ref{sec:Conclusion}.

\label{sec:Conclusion} We have investigated the expectation value in the electron-positron flux from a statistical ensemble of astrophysical sources as well as the uncertainty introduced by our ignorance of the distances and ages of individual sources. We have considered the transport of GCR electrons and positrons in the usual diffusive setup with the same power law injection spectrum for all sources. Their average spatial distribution is modelled combining the distribution of SNRs in galacto-centric radius with a model of the spiral structure. For this as for any other source distribution, the expectation value can be calculated as a sum of power laws in energy where the coefficients are given by the Taylor coefficients of the source distance distribution. We found that the probability density $f_G$ of the Green's function $G$ has a long power law tail. Consequently, the standard deviation is not defined and the probability density for the flux is not a Gaussian but a more general stable distribution. We have quantified the uncertainty around the expectation value by the quantiles of the stable distribution and have also calculated the most-likely flux. We find that the level of fluctuations between different members of the ensemble grows with energy as expected and that most fluxes exhibit a propagation cut-off at a few TeV. In the case without source cut-off the uncertainty interval is in agreement with data from Fermi-LAT and marginally consistent with the measurements by HESS. This can be improved by invoking a source cut-off, e.g. at $E_{\text{cut}} = 20 \, \text{TeV}$. The analytical formulae we have provided allow to consider different types of astrophysical sources with arbitrary spatial distribution functions. We emphasise that the uncertainty in the flux is inherent to the propagation of GCR electrons and positrons and our ignorance of the positions and ages of individual sources. As such it needs to be considered when predicting their fluxes from astrophysical sources.

1012.0805

1012

1012.2388_arXiv.txt

We examine the star formation properties of group and field galaxies in two surveys, SDSS (at z $\sim$ 0.08) and GEEC (at z $\sim$ 0.4). Using UV imaging from the \galex space telescope, along with optical and, for GEEC, near infrared photometry, we compare the observed spectral energy distributions to large suites of stellar population synthesis models. This allows us to accurately determine star formation rates and stellar masses. We find that star forming galaxies of all environments undergo a systematic lowering of their star formation rate between z=0.4 and z=0.08 regardless of mass. Nonetheless, the fraction of passive galaxies is higher in groups than the field at both redshifts. Moreover, the difference between the group and field grows with time and is mass-dependent, in the sense the the difference is larger at low masses. However, the star formation properties of star forming galaxies, as measured by their average specific star formation rates, are consistent within the errors in the group and field environment at fixed redshift. The evolution of passive fraction in groups between z=0.4 and z=0 is consistent with a simple accretion model, in which galaxies are environmentally affected 3 Gyrs after falling into a $\sim 10^{13}$ \Mdotspace group. This long timescale appears to be inconsistent with the need to transform galaxies quickly enough to ensure that star forming galaxies appear similar in both the group and field, as observed.

The star formation history of a galaxy is a function of, at least, stellar mass, redshift and environment. In the local universe, a higher fraction of low mass galaxies are actively forming stars than more massive galaxies \citep[][B04]{kauffmann_mass, brinchmann_sfr}. It has been known for some time that the star formation density of the universe has decreased by at least a factor of 10 in the last 8 or 10 Gyrs \citep{lilly96, madau96, hopkins2004}. There is growing evidence that this reduction with time is seen at all stellar masses \citep{gilbank_roles}. Finally, at least in the local universe, the fraction of star forming galaxies in groups and clusters at fixed stellar mass is lower than the general field \citep{kauffmann_envt, Kimm+09}. Untangling why and to what extent each of stellar mass, redshift and environment determine a galaxy's properties is a fundamental goal of galaxy formation and evolution research. Ultimately, we hope to uncover the physical mechanisms responsible for each correlation. The role of environment has been studied extensively by many groups \citep[for recent reviews, see][]{boselli_review, blanton_review}. Importantly, \citet{weinmann} showed, by separating galaxies based on colour and specific star formation rate (based on H$\alpha$), that the fraction of blue, star forming galaxies decreases with increasing halo mass at fixed luminosity. More recently, using UV derived star formation rates and the same group catalogue, \citet{Kimm+09} finds that the fraction of passive satellites increases with halo mass. These studies were largely focused on the fraction of passive galaxies, rather than the actual star formation rates. Surprisingly, however, there is evidence that galaxies that are forming stars in groups have similar properties to those in the field; it is just that the fraction of those galaxies varies. For example, many authors \citep[eg.][]{strateva,baldry} have found that the local galaxy distribution is bimodal in colour, having a red and a blue peak. \citet{balogh_bimodal} finds that the peak of the red and blue galaxies change relative heights with environment at fixed luminosity. However, importantly, they find no significant difference in the location of the blue peak with environment \citep[but see ][]{wilman_multi}. Recently, \citet{Peng+10} found that the relationship between star formation rate and stellar mass was the same in the highest and lowest density quartile of galaxies. This implies the mechanism that transforms galaxies in dense environments must be rapid. Strangulation, the process in which the hot gas halo surrounding a galaxy is stripped when it becomes a satellite in a large dark matter halo, is often thought to act over timescales of $>$ 2 Gyrs, seemingly in contradiction with the observational need for a rapid timescale \citep{mccarthy}. Therefore, a process that involves the ram pressure stripping of a galaxy's cold gas seems more viable. However, after correcting for the finite number of member galaxies, \citet{Balogh+10} found that the intrinsic scatter between the red fractions in individual galaxy groups and clusters is remarkably small. When directly compared to models for how galaxies are accreted into groups and subsequently into clusters, this small scatter suggests that star formation must be truncated in haloes with mass $\leq$ 10$^{13}$ \Mdoth \citep{mcgee_accretion}. The efficiency of ram pressure stripping of cold gas in such low mass haloes is likely to be poor. The next step is to look at the evolution of the key observational properties for further clues about the relevant mechanisms. It is difficult to obtain a large collection of unbiased and well sampled galaxy groups at high and intermediate redshift and thus the majority of the previous work has been based in the local Universe. To combat this, our collaboration, the Group Environment Evolution Collaboration (GEEC), has undertaken a detailed, multi-wavelength study of galaxy groups at intermediate redshift (0.3 $<$ z $<$ 0.55) \citep{Wilman1}. We have shown that these are truly galaxy groups --- rather than clusters --- as the group-sized velocity dispersions agree with stacked weak lensing and X-ray luminosities \citep{parker, Finoguenov+09}. We have shown that the morphology-environment relation, using either visual or quantitative morphologies, while in place at z=0.4, grows stronger to z=0 \citep{mcgee, wilmanS0}. In addition, as in the local universe, the fraction of [OII] emitting galaxies, infrared excess galaxies or blue galaxies as a function of stellar mass is higher in the field than in the groups \citep{Balogh_smass, wilmanIRAC, Balogh_cnoccol}. [OII] emission can be effectively corrected to be a useful tracer of average star formation rates for large samples \citep{moustakas, gilbank_sdss}. However, it is not clear that these corrections are effective for subsamples of galaxies, such as those that are affected by dense environments \citep{yan_oii, lemaux_oii}. In addition, to properly separate star forming from non-star forming galaxies, it is necessary to probe low star formation rates. It is difficult to attain this level of sensitivity with [OII], especially in low signal to noise spectra. In this paper, we use SED-fit star formation rates, that are driven by UV data from \galex, and stellar masses, that are driven by $K$ band data. This allows us to probe how star formation evolves as a function of environment and stellar mass since z=0.5. In \textsection \ref{sec-data}, we explain the two distinct surveys (GEEC and SDSS) used in the paper as well as the wide array of photometric and spectroscopic data used in both. We will also detail the new \galex (\textsection \ref{sec-galex}) and CFHT Wircam K band (\textsection \ref{sec-kband}) observations that will be used to fit detailed spectral energy distributions. In \textsection \ref {sec-SED}, we explain the SED fits used in this paper including the fitting methodology and the sample of comparison stellar populations that are used to derive physical parameters. Finally, we simulate a sample of galaxies that allow tests of the robustness and accuracy with which we recover physical parameters (\textsection \ref{sec-simlgals}). In \textsection \ref{sec-results}, we examine the environmental dependence of the SSFR - \Mstellar diagram at both redshift epochs. In \textsection \ref{sec-discussion} we discuss the results and derive a plausible toy model for the truncation of star formation in group galaxies. Finally, in \textsection \ref{sec-conclusions}, we discuss our conclusions. In Appendix \ref{sec-physical}, we make direct comparisons between our method of determining physical parameters and other methods from the literature. Throughout this paper, we adopt a \LCDM cosmology with the parameters; $\Omega_{\rm m} = 0.3$, $\Omega_{\Lambda}=0.7$ and $h=H_0/(100 \kmsmpc)=0.75$. Also, in this paper all magnitudes are stated within the AB magnitude system \citep{Oke+83}.

We have fit spectral energy distributions to galaxies in two surveys, SDSS and GEEC. These SEDs use high quality, space based ultraviolet imaging along with optical, and near infrared for GEEC, photometry. We have compared this photometry to large suites of stellar population synthesis models to determine star formation rates and stellar masses. This method nicely reproduced alternative methods of measuring both star formation rates and stellar masses. By examining the results, we conclude the following. \begin{itemize} \item Star forming galaxies of all environments undergo a systematic lowering of their star formation rate between z=0.4 and z=0.08 regardless of mass. \item The star formation properties of star forming galaxies, as measured by their average specific star formation rates, are the same in the group and field environment at fixed redshift. \item The fraction of passive galaxies is higher in groups than the field at both redshifts. However, the difference between the group and field grows with time and is mass dependent, in the sense the the difference is larger at low masses. \item Low mass galaxies at z=0 have group and field passive fractions that can be explained if passive galaxies only exist in groups. \item The evolution of passive fractions in groups between z=0.4 and z=0 is consistent with an accretion model in which galaxies are environmentally affected 3 Gyrs after falling into a 10$^{13}$ \Mdotspace halo/group. \end{itemize} These results present a consistent picture of environmental effects when taken along with our measurements of quantitative morphology in \citet{mcgee}. In that paper, we showed that the fraction of disk galaxies is higher in the field at both redshift and the difference grows larger with time. Also, we found that there was no indication that the disk scaling relations were different in the field or groups. These results all suggest that only a fraction of galaxies in groups must be actively truncated at any given time. If the mechanism is a quick one, like ram pressure stripping, then the truncation time might be short enough that only a small fraction of group galaxies are affected at a given time. This would allow the bulk of the star forming galaxies to remain unchanged but still allow the observed evolution. However, ram pressure stripping is likely not effective in galaxy groups. Based on the timescale suggested by our accretion model, strangulation seems like a suitable candidate for environmental mechanisms. However, it is unclear if strangulation can allow galaxies to remain apparently unaffected for some time, thereby appearing to act only on a fraction of group galaxies at a time. In semi-analytic models, strangulation produces too many 'green' galaxies, which would likely alter the disk and star formation properties \citep{font, Balogh_cnoccol}. These results suggest that further constraints can be applied in two ways. First, the detailed study of individual galaxy groups and the orbits of their constituent galaxies within them can determine if ram pressure stripping of the cold gas is a viable mechanism. Secondly, the continued hunt for elusive `green' transition galaxies, perhaps at high redshift where the accretion rates of galaxies into groups is higher \citep{mcgee_accretion, geec_hiz}, will determine the viability of gentle, strangulation like mechanisms.

1012.2388

1012

1012.5757_arXiv.txt

The intrinsic shapes of elliptical galaxies and disks have been extensively studied in the literature. However, bulges appear to be less studied, even if they account for about $25\%$ of the stellar mass of the local universe \cite{driver07}. The study of the intrinsic shape of bulges presents similarities, advantages, and drawbacks with respect to that of elliptical galaxies. For bulges, the problem is complicated by the presence of other luminous components and requires the accurate isolation of their light distribution. On the other hand, the presence of the galactic disk allows for the accurate constraining of the inclination of the bulge under the assumption that the two components share the same polar axis. Although the kinematical properties of many bulges are well described by dynamical models of oblate ellipsoids which are flattened by rotation with little or no anisotropy \cite{daviesillingworth83, corsini99, pignatelli01}, the twisting of the bulge isophotes \cite{zaritskylo86} and the misalignment between the major axes of the bulge and disk \cite{bertola91, mendezabreu08} observed in several galaxies are not possible if the bulge and disk are both axisymmetric. These features were interpreted as the signature of bulge triaxiality. This idea is also supported by the presence of non-circular gas motions \cite{gerhardvietri86, falconbarroso06, pizzella08} and a velocity gradient along the galaxy minor axis \cite{corsini03, coccato04, coccato05}. Perfect axisymmetry is also ruled out when the intrinsic shape of bulges is determined by statistical analyses based on their observed ellipticities. Bertola et al. \cite{bertola91} measured the bulge ellipticity and the misalignment between the major axes of the bulge and disk in 32 S0--Sb galaxies. They found that these bulges are triaxial with mean axial ratios $\langle B/A \rangle=0.86$ and $\langle C/A \rangle=0.65$. In contrast, measurements of $\langle B/A \rangle=0.79$ for the bulges of 35 early-type disk galaxies and $\langle B/A \rangle=0.71$ for the bulges of 35 late-type spirals were found by Fathi \& Peletier \cite{fathipeletier03}. None of the 21 disk galaxies with morphological types between S0 and Sab studied by Noordermeer \& van der Hulst \cite{noordermeervanderhulst07} harbors a truly spherical bulge. They obtain a mean flattening $\langle C/A \rangle=0.55$. Mosenkov et al. \cite{mosenkov10} obtained a median value of the flattening $\langle C/A \rangle=0.63$ for a sample of both early and late-type edge-on galaxies in the near infrared. In M\'endez-Abreu et al. \cite{mendezabreu08}(Paper I) we measured the structural parameters of a sample of 148 unbarred early-to-intermediate spiral galaxies using the GASP2D algorithm to analyze their near-infrared surface-brightness distribution. The probability distribution function (PDF) of the bulge equatorial ellipticity was derived from the distributions of observed ellipticities of bulges and misalignments between bulges and disks. We proved that about $80\%$ of the sample bulges are not oblate but triaxial ellipsoids with a mean axial ratio $\langle B/A\rangle = 0.85$. In this work, (see M\'endez-Abreu et al. \cite{mendezabreu10} for details) we introduce a new method to derive the intrinsic shape of bulges under the assumption of triaxiality. This statistical analysis is based upon the analytical relations between the observed and intrinsic shapes of bulges and their surrounding disks and it is applied to the galaxy sample described in Paper I. The method make use only of photometric data and have been conceived to be completely independent of the studied class of objects, and it can be applied whenever triaxial ellipsoids embedded in (or embedding) an axisymmetric component are considered.

\label{sec:conclu} In this work, we have developed a new method to derive the intrinsic shape of bulges. It is based upon the geometrical relationships between the observed and intrinsic shapes of bulges and their surrounding disks. We assumed that bulges are triaxial ellipsoids with semi-axes of length $A$ and $B$ in the equatorial plane and $C$ along the polar axis. The bulge shares the same center and polar axis of its disk, which is circular and lies on the equatorial plane of the bulge. The intrinsic shape of the bulge is recovered from photometric data only. They include the lengths $a$ and $b$ of the two semi-major axes of the ellipse, corresponding to the two-dimensional projection of the bulge, the twist angle $\delta$ between the bulge major axis and the galaxy line of nodes, and the galaxy inclination $\theta$. The method is completely independent of the studied class of objects, and it can be applied whenever a triaxial ellipsoid embedded in (or embedding) an axisymmetric component is considered. We analyzed the magnitude-limited sample of 148 unbarred S0--Sb galaxies, for which we have derived (Paper I) their structural parameters by a detailed photometric decomposition of their near-infrared surface-brightness distribution. We derived the triaxiality parameter, as defined by \cite{franx91}, for all of them. We found that it follows a bimodal distribution with a minimum at $T=0.55$ and two maxima at $T=0.05$ (corresponding to oblate axisymmetric or nearly axisymmetric ellipsoids) and $T=0.85$ (strongly prolate triaxial ellipsoids), respectively. This bimodality is driven by bulges with S\'ersic index $n > 2$ or alternatively by bulges of galaxies with a bulge-to-total ratio $B/T > 0.3$. Bulges with $n \leq 2$ and bulges of galaxies with $B/T \leq 0.3$ follow a similar distribution, which is different from that of bulges with $n > 2$ and bulges of galaxies with $B/T > 0.3$. The different distribution of the intrinsic shapes of bulges according to their S\'ersic index gives further support to the presence of two bulge populations with different structural properties: the classical bulges, which are characterized by $n > 2$ and are similar to low-luminosity elliptical galaxies, and pseudobulges, with $n \leq 2$ and characterized by disk-like properties. The correlation between the intrinsic shape of bulges with $n \leq 2$ and those in galaxies with $B/T \leq 0.3$ and between bulges with $n > 2$ and those in galaxies with $B/T > 0.3$ agrees with the correlation between the bulge S\'ersic index and bulge-to-total ratio of the host galaxy, as recently found by \cite{droryfisher07} and \cite{fisherdrory08}. The observed bimodal distribution of the triaxiality parameter can be compared to the properties predicted by numerical simulations. \cite{cox06} studied the structure of spheroidal remnants formed from major dissipationless and dissipational mergers of disk galaxies. Dissipationless remnants are triaxial with a tendency to be more prolate, whereas dissipational remnants are triaxial and tend be much closer to oblate. In addition, \cite{hopkins10} used semi-empirical models to predict galaxy merger rates and contributions to bulge growth as functions of merger mass, redshift, and mass ratio. They found that high $B/T$ systems tend to form in major mergers, whereas low $B/T$ systems tend to form from minor mergers. In this framework, bulges with $n \leq 2$, which shows a high fraction of oblate axisymmetric (or nearly axisymmetric) shapes and have $B/T \leq 0.3$, could be the result of dissipational minor mergers. A more complex scenario including both major dissipational and dissipationless mergers is required to explain the variety of intrinsic shapes found for bulges with $n > 2$ and $B/T > 0.3$. However, high-resolution numerical simulations in a cosmologically motivated framework that resolves the bulge structure are still lacking. The comparison of a larger sample of bulges with a measured intrinsic shape and covering the entire Hubble sequence with these numerical experiments is the next logical step in addressing the issue of bulge formation.

1012.5757

1012

1012.0344_arXiv.txt

We report observations of a white-light solar flare ({\tt SOL2010-06-12T00:57}, M2.0) observed by the \textit{Helioseismic Magnetic Imager} (HMI) on the \textit{Solar Dynamics Observatory} (SDO) and the \textit{Reuven Ramaty High-Energy Solar Spectroscopic Imager} (RHESSI). The HMI data give us the first space-based high-resolution imaging spectroscopy of a white-light flare, including continuum, Dop\-pler, and magnetic signatures for the photospheric Fe~{\sc i} line at 6173.34~\AA~and its neighboring continuum. In the impulsive phase of the flare, a bright white-light kernel appears in each of the two magnetic footpoints. When the flare occurred, the spectral coverage of the HMI filtergrams (six equidistant samples spanning $\pm$172~m\AA~around nominal line center) encompassed the line core and the blue continuum sufficiently far from the core to eliminate significant Doppler crosstalk in the latter, which is otherwise a possibility for the extreme conditions in a white-light flare. RHESSI obtained complete hard X-ray and $\gamma$-ray spectra (this was the first $\gamma$-ray flare of Cycle~24). The Fe~{\sc i} line appears to be shifted to the blue during the flare but does not go into emission; the contrast is nearly constant across the line profile. We did not detect a seismic wave from this event. The HMI data suggest stepwise changes of the line-of-sight magnetic field in the white-light footpoints.

Solar flares are explosive phenomena visible in all regions of the solar atmosphere, and were originally discovered by \inlinecite{1859MNRAs..20...13C} via emission in white light. Since then, many other manifestations of flares have been discovered in the outer atmosphere and heliosphere, including signatures of acoustic waves penetrating into the solar interior. According to general consensus, the flare phenomenon corresponds to the sudden release of energy stored in the corona via the slow buildup of excess magnetic energy, which ultimately originated via dynamo action within the convective envelope. Many of the mechanisms remain ill-understood, including the nature of the initial plasma instability that sets the flare off. The detection in white light immediately implies a large concentration of the released energy in the lower atmosphere. In the upper atmosphere and corona, X-ray and $\gamma$-ray signatures show that the energy release has the property of strong particle acceleration, to the extent that major fractions of the total energy appear to be in electrons above 10~keV and protons above 1~MeV. These high-energy radiations, and the acceleration of an associated coronal mass ejection (CME), define the ``impulsive phase'' of a flare; other related energy release may take the form of gentler heating. The original Carrington observation still pre\-sents several open questions. The white-light flare remains the energetically decisive flare observational signature because most of the flare energy is in the visible and near-UV \cite{2006JGRA..11110S14W,2007ApJ...656.1187F}. We now think that the visible continuum is enhanced in all flares, but that for the weaker ones the signal is lost in the spatial and temporal brightness fluctuations of the photosphere. The continuum emission appears in the impulsive phase and is located in the deep solar atmosphere, even apparently reaching the ``opacity minimum'' region of the spectrum near 1.56~$\mu$m \cite{2004ApJ...607L.131X}. In spite of this supposedly photospheric signal, strong evidence also implicates the chromosphere, since the continuum emission includes clear signatures of recombination radiation \cite{1986lasf.conf..142N,2010arXiv1001.1005H}. A part of this evidence is the strong association of the white-light continuum with hard X-rays \cite{1970SoPh...13..471S,1975SoPh...40..141R,1992PASJ...44L..77H}. A 10-keV electron has a limited range in a plasma, and cannot reach the photosphere if accelerated in the corona as in the standard thick-target model \cite{1971SoPh...18..489B,1972SoPh...24..414H}. Accordingly the elucidation of the paths of energy propagation from coronal magnetic storage to the lower atmosphere has great importance. The Poynting flux could replace the electron beams of the heretofore-standard thick-target model for this purpose \cite{2008ApJ...675.1645F,2009ApJ...695.1151B}, but in any case the strong acceleration of non-thermal electrons above 10~keV remains a requirement. Observations from the \textit{Helioseismic and Magnetic Imager} (HMI) onboard the {\it Solar Dynamics Observatory} (SDO) give us the first true imaging spectroscopy of flare effects in the photosphere at high spectral and spatial resolution \cite{2010AAS...21630801S} from space. Previous ground-based observations, typically with slit spectrographs and film readout, have not provided such comprehensive coverage (see \citeauthor{1989SoPh..121..261N} \citeyear{1989SoPh..121..261N} and \citeauthor{2007BCrAO.103...63B} \citeyear{2007BCrAO.103...63B}, for a discussion of this limited material and the conclusions drawn from it). The new data clearly resolve the profile of the Fe~{\sc i}~line, in each $\approx\ $0.5$''$ pixel and 45-second time step. The HMI data and RHESSI data for this flare confirm the intimate relationship between flare effects in the lower atmosphere, and high-energy processes revealed by hard X-ray and $\gamma$-ray emissions. We exploit the new features of HMI to characterize the continuum emission and line-of-sight Doppler and magnetic properties of the two magnetic footpoints that mark a small but exceptionally impulsive white-light flare. This study is intended as preparation for future flare observations by SDO and other space-borne and ground-based facilities in Cycle 24. \begin{table*} \caption{Flare timeline in HMI data (45-second data frames)} \centering \smallskip \begin{tabular}{l l l l l} Interval name & Start (UTC) & HMI continuum & HMI Doppler & RHESSI 100~keV\\ \hline Preflare & 00:54:11 & no excess & no Doppler & no detection \\ Brightening & 00:55:41 & $>$10\% increase & 2~km s$^{-1}$ blue & bright \\ Postflare & 01:00:11 & no excess & complex & no detection \\ \hline \end{tabular} \end{table*} \label{tab:signs}

The HMI imaging spectroscopy of this flare resulted in maps of the line profile of the photospheric line Fe~{\sc i} at 6173.34\AA. The information includes line width, depth, Doppler shift, and line-of-sight magnetic field via observations in two polarizations. They also record the neighboring continuum near 6173.34~\AA. This capability greatly improves our understanding of the lower solar atmosphere during flares, since the full information of the line profile can be interpreted in terms of the physical conditions there. The full imaging spectroscopy allows us to identify the continuum brightenings with the footpoint regions of coronal loop structures via reference to AIA images. A fuller analysis of these data (and those of EVE) is outside the scope of this paper. For the flare studied here we find that the line shifts in wavelength, but remains in absorption. The intensity in the core of the line appears to have approximately the same flare excess as the nearby continuum (cf. \opencite{2010ApJ...722.1514P}). We also find strong blueshifts of the line at both footpoints. Both of these findings are significant of multiple possibilities. Backwarming \cite{1989SoPh..124..303M} due to the observed continuum should raise the photospheric temperature, which could, in principle, change both the strength and width of the line. The observation of numerous other white-light flares with HMI should restrict these possibilities; we cannot at present rule out the possibility that the blueshifts could be an artifact of the sampling sequence. The transient blueshift could signify a photospheric medium moving toward the spacecraft, shifting the absorption line accordingly. Alternatively, it could be the result of red-shifted component of line emission, from a down-flowing heated chromosphere\footnote{The latter would be consistent with red-shifted H$\alpha$ emission in reaction to chromospheric ablation at higher altitudes \cite{2009ApJ...699..968M} and in Na\,D$_1$-line emission seen by \inlinecite{2005ApJ...630.1168D}. The chromospheric emission would have to be insufficiently strong to drive the photospheric line into emission, but the resultant of the superposition would be a blue-shifted absorption profile.}. Detailed radiative-transfer modeling of appropriate scenarios is needed to address this question. It is also tantalizing is that the transient blue shift is so apparently at odds with the transient red (or mixed) Doppler shifts seen in other observations \cite{koso2006K,2007AdSpR..40.1921B}. For a magnetic region as far limbward as AR11081 at the time of {\tt SOL-2010-06-12T00:57}, it is important to consider horizontal motion of the medium. This could either be motion driven by a magnetic jerk or magnetic deflection of motion that would have been vertical except for a strong, inclined magnetic field. A basic control question for this hypothesis, then, is whether flares from near disk center show the transient red shifts generally seen in MDI observations, as Doppler observations of these should be insensitive to horizontal motion. A final consideration, given the suddenness and short duration of the HXR profile in {\tt SOL2010-06-12T00:57}, is the possibility of a blueshift artificially caused by aliasing. The time separation of red and blue spectroheliograms, in the presence of a varying line intensity, could result in a spurious Doppler signal. This possibility can be controlled by comparisons with GONG observations of {\tt SOL2010-06-12-T00:57}, in which temporal aliasing is greatly reduced in integrations of respective intensity and Doppler signatures for the full duration represented by the record time. This comparison is being undertaken in a study in progress that benefits from new analysis techniques that allow us to compensate for noise introduced by variations in atmospheric seeing quality \cite{2008SoPh..251..627L}. The HMI observations are consistent with the idea that the flare emission at this wavelength has a substantial component of Paschen continuum from hydrogen recombination at higher altitudes, a conclusion also consistent with observations of the Balmer and even Paschen continuum edges in the spectra of other white-light flares \cite{1989SoPh..121..261N}. To understand these results quantitatively will require modeling beyond the scope of this paper, and of course to draw any general conclusion would require the observation of other flares in this manner as well. The HMI observations clearly point to the need for higher temporal resolution along with good spatial resolution in white light. Ground-based observations (\textit{e.g.,} \opencite{2008ApJ...688L.119J}) could extend this work, but we also expect many more interesting flare observations from the powerful instrumentation on SDO. We expect that comparable HMI data will become available for many other flares and that the modeling of the lower solar atmosphere will solve several outstanding problems of interpretation, including the nature of the impulses responsible for sunquakes \cite{1998Natur.393..317K}. \bigskip\noindent{\bf Acknowledgments:} The Berkeley group was supported by NASA under contract NAS 5-98033 for RHESSI, and the SDO/HMI by contract NAS5-02139 to Stanford University. We thank the other members of the SDO/HMI team for special help during a very busy time, and for building such a fine instrument. These data have been used courtesy of NASA/SDO and the AIA, EVE, and HMI science teams. The author list of this paper consists of those persons actively involved in actually writing the paper, rather than making it possible, and after the first author's name the order has been randomized. \appendix \label{sec:appendix} The HMI Dopplergrams and intensities used for standard helioseismic applications of the data consist of a spatial-temporal interpolation, using 12~filtergrams at six different wavelengths and two polarization states, for each observable. These interpolations, in conjunction with the scanning of the Fe~{\sc i} line, are devised to minimize the effects of aliasing in \textit{p}-mode recognition. The standard interpolation for this purpose includes contributions from 135 seconds before and after the time assigned to a Dopplergram, with coefficients that are negative in the ranges $\pm$45\,--\,90~seconds (see Figure~\ref{fig:filter}). Because of this, the response of the standard Dopplergrams to a sufficiently sharp white-light flare can be an apparent reduction in intensity preceding the flare, \textit{i.e.}, an apparent ``black-light flare'' preceding the white-light excess (see Figure~\ref{fig:artifact}). \begin{figure}[htb] \centering \includegraphics[angle=90,width=0.5\columnwidth]{hugh.eps} \caption{Plot of the function specifying the interpolation coefficients applied to filtergrams for computation of Doppler, intensity, and line-of-sight magnetic maps in HMI time series. Note the up-to-15\% reversal in the value of the function in the intervals $\pm$45--90~seconds.} \label{fig:filter} \end{figure} \begin{figure} \centering \includegraphics[width=\columnwidth]{FIGURE1.eps} \caption{Intensity artifact prior to a white-light flare in helioseismically interpolated intensity maps. Upper left frame shows pre-flare intensity. Upper right frame shows the pre-flare line-of-sight magnetic field. Bottom row shows consecutive intensity-difference images. Bottom left shows the pre-flare difference; bottom right shows the white-light-flare difference. Note the apparent reduction in intensity in the flare footpoints in the lower left frame.} \label{fig:artifact} \end{figure}

1012.0344

1012

1012.0172_arXiv.txt

s{ We study the dynamical evolution of cosmological models with the Robertson-Walker symmetry with a scalar field non-minimally coupled to gravity and barotropic matter. For this aim we use dynamical system methods. We have found a type of evolutional path which links between all important events during the evolution, the cosmological singularity of finite time, inflation, radiation and matter dominating epoch and the accelerated phase expansion of the universe. We point out importance of finding the new generic solution called a twister solution for a deeper description of the evolution of the Universe. We demonstrate that including the non-minimal coupling leads to a new, richer evolutional cosmological scenario in comparison to the case of minimal coupling.}

The standard method of description matter content in cosmology bases on the concept of perfect fluid approximation, where pressure and energy density satisfy the equation of state. Different cosmological epochs constitute solutions of the Einstein equations corresponding different forms of the equation of state, usually postulated in a linear form with respect to the energy density. On the other hand if we study very early stages evolution of the Universe then characterization of matter content in terms of the barotropic equation of state is not adequate. In the quantum epoch including matter in the form of scalar field with the potential seems to be more suitable. We propose to describe matter in the form of both barotropic matter and single scalar field with the potential. We also assumed that there is present a nonzero coupling constant scalar field to the gravity. A non-minimal coupling appeared naturally in quantum theory of the scalar field as generated by quantum corrections or required by renormalization of the theory.\cite{Faraoni:1996rf} The value of this coupling constant becomes important for cosmology. Recently, it has been constructed an extended model of inflation with a non-minimal coupling between the inflaton field and the Ricci scalar curvature.\cite{Bezrukov:2007ep,DeSimone:2008ei,Barvinsky:2008ia,Pallis:2010wt} It was shown in particular that the non-minimal inflation can be realized within the Standard Model (SM) or minimal extension of it.\cite{Lerner:2009xg} In the SM the Higgs boson inflation and dark matter is considered by Clark et al.\cite{Clark:2009dc} The non-minimal coupling plays also an important role in the context of description of quintessence epoch.\cite{Szydlowski:2008in,Hrycyna:2009zf,Hrycyna:2010yv} The main aim of this paper is to present a generic solution of the non-minimal coupling cosmology called the twister solution because of the shape of its trajectory in the phase space. We explore here methods of dynamical systems because their advantage of representing all solutions for all admissible initial conditions. The phase space is organized by critical points which represent stationary solutions for which right-hand sides of the dynamical system vanish and trajectories joining them represent the evolution of the system. A new type of evolution appears only if the coupling constant is different from minimal and conformal. This solution leads naturally to the quintessence epoch through the twister solution. We characterize properties of this acceleration by calculation of so called state-finder parameters. We assume the spatially flat FRW universe filled with the non-minimally coupled scalar field and barotropic fluid with the equation of the state coefficient $w_{m}$. The action is \begin{equation} S = \frac{1}{2}\int \ud^{4}x \sqrt{-g} \left[\frac{1}{\kappa^{2}}R - \ve \Big(g^{\mu\nu}\partial_{\mu}\phi\partial_{\nu}\phi + \xi R \phi^{2}\Big) - 2U(\phi) \right] + S_{m}, \end{equation} where $\kappa^{2}=8\pi G$, $\ve = +1,-1$ corresponds to canonical and phantom scalar fields, respectively, the metric signature is $(-,+,+,+)$, $R=6\left(\frac{\ddot{a}}{a}+\frac{\dot{a}^{2}}{a^{2}}\right)$ is the Ricci scalar, $a$ is the scale factor and a dot denotes differentiation with respect to the cosmological time $t$ and $U(\phi)$ is the scalar field potential function. $S_{m}$ is the action for the barotropic matter part. The dynamical equation for the scalar field we can obtain from the variation $\delta S/\delta \phi = 0$ \begin{equation} \ddot{\phi} + 3 H \dot{\phi} + \xi R \phi + \ve U'(\phi) =0, \end{equation} and energy conservation condition from the variation $\delta S/\delta g^{\mu\nu}=0$ \begin{equation} \mathcal{E}= \ve \frac{1}{2}\dot{\phi}^{2} + \ve3\xi H^{2}\phi^{2} + \ve3\xi H (\phi^{2})\dot{} + U(\phi) + \rho_{m} - \frac{3}{\kappa^{2}}H^{2}. \end{equation} Then the Einstein equation for the flat FRW model and the conservation condition read \begin{equation} \frac{3}{\kappa^{2}}H^{2} = \rho_{\phi} + \rho_{m}, \quad \textrm{and} \quad \dot{H} = -\frac{\kappa^{2}}{2}\Big[(\rho_{\phi}+p_{\phi}) + \rho_{m}(1+w_{m})\Big] \end{equation} where the energy density and the pressure of the scalar field are \begin{eqnarray} \rho_{\phi} & = & \ve\frac{1}{2}\dot{\phi}^{2}+U(\phi)+\ve3\xi H^{2}\phi^{2} + \ve 3\xi H (\phi^{2})\dot{},\\ p_{\phi} & = & \ve\frac{1}{2}(1-4\xi)\dot{\phi}^{2} - U(\phi) + \ve\xi H(\phi^{2})\dot{} - \ve2\xi(1-6\xi)\dot{H}\phi^{2} - % \ve3\xi(1-8\xi)H^{2}\phi^{2} + 2\xi\phi U'(\phi). \end{eqnarray}

In this paper we pointed out the presence of the new interesting solution for the non-minimally coupled scalar field cosmology which we called the twister solution (because of the shape of the corresponding trajectory in the phase space). This type of the solution is very interesting because in the phase space it represents the 3-dimensional trajectory which interpolates different stages of evolution of the universe, namely, the radiation dominated, dust filled and accelerating universe. We are able to find linearized solutions around all these intermediate phases, and hence, parameterizations for $w_{\rm{eff}}(a)$ in different epochs of the universe history. It is interesting that the presented structure of the phase space is allowed only for non-zero value of the coupling constant, therefore it is a specific feature of the non-minimally coupled scalar field cosmology. Our general conclusion is that in the description of dynamical complexity the cosmic evolution including both barotropic matter and non-minimally scalar field leads to a new richer dynamics.

1012.0172

1012

1012.4421_arXiv.txt

The potential of combining Adaptive Optics (AO) and Lucky Imaging (LI) to achieve high precision astrometry and differential photometry in the optical is investigated by conducting observations of the close 0\farcs1 brown dwarf binary GJ569Bab. We took 50000 $I$-band images with our LI instrument FastCam attached to NAOMI, the 4.2-m William Herschel Telescope (WHT) AO facility. In order to extract the most of the astrometry and photometry of the GJ569Bab system we have resorted to a PSF fitting technique using the primary star GJ569A as a suitable PSF reference which exhibits an $I$-band magnitude of $7.78\pm0.03$. The AO+LI observations at WHT were able to resolve the binary system GJ569Bab located at $4\farcs 92 \pm 0\farcs05$ from GJ569A. We measure a separation of $98.4 \pm 1.1$ mas and $I$-band magnitudes of $13.86 \pm 0.03$ and $14.48 \pm 0.03$ and $I-J$ colors of 2.72$\pm$0.08 and 2.83$\pm$0.08 for the Ba and Bb components, respectively. Our study rules out the presence of any other companion to GJ569A down to magnitude I$\sim$ 17 at distances larger than 1\arcsec. The $I-J$ colors measured are consistent with M8.5-M9 spectral types for the Ba and Bb components. The available dynamical, photometric and spectroscopic data are consistent with a binary system with Ba being slightly (10-20\%) more massive than Bb. We obtain new orbital parameters which are in good agreement with those in the literature.

The lucky imaging (LI) technique proposed by \citet{FriedD:probgl} attracted the attention by professional astronomers once low read-out noise detectors became available \citep[e.g.][]{BaldwinJ:diffl8, TubbsR:difflc, LawN:luckih}. Recently it has been realized that the combination of LI and Adaptive Optics (AO) can benefit mutually and provide high-resolution imaging close to the diffraction limit at optical and near-infrared (NIR) wavelengths \citep[e.g. see][]{GladyszS:luckis,LawN:GettLAO,KervellaP:cloce}. This can be obtained by using a larger fraction of data in the LI selection and/or keeping images with a better Strehl ratio than in conventional LI observations. In this paper we present the results of combining the LI and AO techniques to produce high-angular resolution and high-contrast imaging in the optical of the multiple system GJ569Bab which is a benchmark in the study and characterization of brown dwarfs (BDs). GJ569A is an M2.5V cromospherically active star lying at a distance of 9.6-9.8 pc \citep{PerrymanM:HIPPARCOS,vanLeeuwenF:hippnr}. \citet{ForrestW:possbd} identified a faint companion to GJ569 and argued the potential brown dwarf nature of such companion. Using Keck AO observations \citet{MartinE:discv} resolved GJ569B as a binary brown dwarf system with a separation of $\sim 0\farcs1$, a total mass of the system in the range 0.09-0.15~\MS, an age in the range 0.12-1.0~Gyr and an orbital period $\sim 3$ yr. Further AO-based observations with the Keck telescope \citep[e.g. ][]{LaneB:orbbd, ZapateroM:dynmbb, SimonM:gl569ms, KonopackyQ:highpd} and the Subaru and HST telescopes \citep{ZapateroM:lithd, MartinE:reshs} have allowed precise determination of the dynamical masses and orbital parameters of the binary system GJ569Bab (see Table~\ref{tab:orbit} in this paper) as well as a precise determination of the spectral types of the GJ569B components: M8.5-9V and M9V for the Ba and Bb components, respectively \citep{LaneB:orbbd,MartinE:reshs}. For BDs an estimate of the mass is essential to determine their properties and evolution. A way to achieve a direct measurement of masses is to observe close binary systems where the short orbital period allows for a complete sampling of the orbit and from here a precise determination of the dynamical masses of the pair. Up to now a complete characterization of the orbital motion has been achieved for a few BD binaries \citep[e.g. see][]{DupuyT:dynams, KonopackyQ:highpd}. In this context GJ569Bab constitutes a unique laboratory where to test the stellar evolutionary models as it is among the shortest known period BD binary system. This has allowed to determine the orbit of the system over several periods and from here a precise determination of its dynamical mass. The high angular resolution requested to spatially resolve faint systems like GJ569B into its components has been so far achieved with 10-m class telescopes and AO in the NIR or with the HST.. Our motivation to perform observations of GJ569 with LI+AO on a 4-m class telescope was twofold. First, to test the potential of this technique for high-angular resolution and high-contrast imaging in the optical regime. Second, to shed light on whether GJ569B is actually a triple system as suggested in some works in the literature \citep{MartinE:discv,KenworthyM:Gl569Bay,SimonM:gl569ms}. In Section~\ref{Sect:Setup} we briefly describe the instrumental set-up. Section~\ref{Sect:Observations} describes the observations. Section~\ref{Sect:Analysis} reports the data calibration, reduction and analysis. Section~\ref{Sect:Results} focuses on the photometry and astrometry of the GJ569 system components and discussion of the results. We provide our conclusions in Sect.~\ref{Sect:Conclusions}.

\label{Sect:Conclusions} We have presented results showing the potential for high precision astrometry, differential photometry and high contrast imaging using a Lucky Imaging instrument coupled to an adaptive-optics system. Our results indicate that with 4-m class telescopes equipped with a moderately low order adaptive optics system it is possible to achieve angular resolutions better than 0\farcs1 in the I-band. This is comparable to what is achieved in the Ks band with the use of AO-systems at 8-10 m class telescopes. This work is part of an effort to determine the feasibility of a Lucky Imaging instrument to be coupled on the future AO system at the GTC telescope (GTCAO) and what the expectations in terms of high contrast imaging and angular resolution should be expected from such a combination. Our work has focused on the observation of the GJ569 system which contains a benchmark brown dwarf binary. The application of our LM PSF fitting technique to the GJ569 image obtained with the WHT, where the GJ569B components are resolved, allows to achieve high precision relative photometry (to a few millimagnitudes) and astrometry (to a few mas) thanks to the availability of meaningful error bars associated to each pixel in the fitted image. The potential for detection of faint companions has been addressed by looking at the $3\sigma$ detectability curves in Fig.~\ref{fig:ContrastWHT}. On the images directly from the frame selection procedure, we distinguish two regions in which the detectability behaves differently versus increasing the percentage of images being employed. In regions where the image is dominated by the wide swallow halo of the primary PSF, the detectability is improved by restricting the percentage of images being employed since in that way there is more energy in the core and less in the halo. At large distances from the primary, the image is dominated by background/detector noise and the detectability is improved by simply adding as many images as possible. On the wavelet-processed images, both in the inner and outer regions, we only see a benefit on the increase of the percentage of images used in the frame selection. With the wavelet-processed images we observe a magnitude gain in the inner region (1.7 magnitudes at 1\arcsec with respect to the non-processed image) but far away from the primary no net gain as those parts of the image are dominated by background/detector noise. We have measured a differential magnitude at I band between GJ569Ba and GJ569Bb $\Delta m_{BaBb}=0.622 \pm 0.017$. When used in conjunction of $\Delta m_{BaBb}$ in the J, H and K bands by previous works, fits well with the spectral determination of M8.5-M9 for the brown dwarf binary derived in the near-infrared by \citet{LaneB:orbbd}. Our results in $I$-band and those in $J$, $H$ and $K$-bands in \citet{LaneB:orbbd} clearly indicate that Ba is brighter than Bb. This together with the $I-J$ color favors a half subspectral class earlier for Ba than for Bb (see Fig.~\ref{fig:ColorIJ}). The astrometric quality achieved with FastCam allows to locate two new points on the GJ569Bb orbit around GJ569Ba. The orbits in the literature and the one derived including our points do not differ significantly and therefore the orbital parameters are in perfect agreement with those previously published, although our mass estimate of 0.116$\pm$0.007~\MS using the updated Hipparcos parallax distance of 9.65$\pm$0.16~pc in \citet{vanLeeuwenF:hippnr}. Our mass estimate of the multiple GJ569B system is somewhat smaller, but within error bars of previously published values except for the newly derived mass of $0.140^{+0.009}_{-0.008}$~\MS in \citet{DupuyT:studpd}. Our WHT data point on June 2009 falls within 1 sigma from the \citet{DupuyT:studpd} orbital solutions and our own orbital solution although our NOT data point on July 2008 is more consistent with previous orbital solutions in \citet{ZapateroM:dynmbb,SimonM:gl569ms,KonopackyQ:highpd}. For a substantial refinement of the orbital parameters it would be necessary to sample the whole orbit with similar uncertainties as those derived from our observation at the WHT on June 2009. The data available on the GJ569B system is consistent with a primary of $0.081 \pm 0.010$~\MS and a secondary of $0.059\pm 0.007$~\MS. If Ba were a binary system then the Bab component would have a mass < 0.018 \MS. A simple qualitative analysis on the deviations from a Keplerian orbit allows to place an upper limit to the product of mass and orbit semi-axis of this object.

1012.4421

1012

1012.1817_arXiv.txt

We conduct a study of K to M type stars to investigate the activity and the rotation limit in the Hyades. We use a sample of 40 stars in this intermediate-age cluster ($\approx$625 Myr) to probe stellar rotation in the threshold region where stellar activity becomes prevalent. Here we present projected equatorial velocities ($v_{rot}\sin i$) and chromospheric activity measurements (H$_\alpha$) that indicate the existence of fast rotators in the Hyades at spectral types where also the fraction of stars with H$_\alpha$ emission shows a rapid increase (``H$_\alpha$ limit''). The locus of enhanced rotation (and activity) thus seems to be shifted to earlier types in contrast to what is seen as the rotation limit in field stars. The relation between activity and rotation appears to be similar to the one observed in fields stars.

Solar-type stars are mostly fast rotators and magnetically active when they are young. Their magnetic fields drive stellar winds, which rotationally slow-down the star by means of angular momentum transfer. The stellar spin-down over time is empirically quantified by the so-called ''Skumanich-law`` as $\omega \propto t^{-\frac{1}{2}}$ \citep{Skumanich1972}. This relation, however, becomes invalid at very low masses. Among the field stars, it is observed that at the transition to fully convective stars at early M-type ($\approx0.3\,M_{\odot}$), the rotational braking efficiency changes, and fast rotation ($v_{\rm{rot}}>3\,$km/s) becomes predominant \citep{Delfosse1998, Mohanty2003, ReinersBasri2008}. The threshold between slow and rapid rotation is thought to be age-dependent \citep{Hawley1999}, so that young cluster stars are expected to show a rotation limit shifted towards higher masses or earlier spectral types, compared to (old) field stars. Stellar rotation and magnetic activity are tightly linked by the underlying dynamo processes. In (young) clusters, it is observed that the fraction of active stars (eg.\ with chromospheric H$_\alpha$ emission) sharply increases at different masses depending on the cluster age \citep[``H$_\alpha$ limit'',][]{Hawley1999}. At younger age, enhanced magnetic activity is seen at higher masses (ie.\ earlier spectral types) than it is for older clusters. However, it is elusive whether this change in the locus of the H$_\alpha$ limit is also due to an increase of the rapid rotation rate \citep{Radick1987, Stauffer1987}. Previous studies have focussed on the evolution of the H$_\alpha$ limit in young and intermediate age clusters \citep{Stauffer1997, Hawley1999, Reid1995} or field stars \citep{West2004}, but rotational velocities have only been measured extensively for field stars all across the main sequence \citep{Delfosse1998, Mohanty2003, West2008, ReinersBasri2008}. However, for open clusters such as the Hyades, $v_{rot}\sin i$ measurements have concentrated on earlier spectral types F to K \citep[eg.][]{Radick1987}, and on the very low-mass regime \citep[M-type and below; eg.][]{ReidMahoney2000}, so that in the mid-K to early M-type range (hence in the range of the H$_\alpha$-limit) rotational velocities are scarce for the Hyades. The present work adresses this scarcity between spectral classes K and M, and probes the coexistence of enhanced activity at the H$_\alpha$ limit and rapid rotation for young stars in the case of the intermediate aged Hyades. We thus aim to determine if rapid rotation occurs at a different threshold in the Hyades with respect to field stars.

Our sample of Hyads shows increased rotation alongside with H$_\alpha$ activity at spectral types $\approx$M0 and later. We see evidence that young, early M- type stars rotate faster than field stars do. This is possibly influenced by the different contraction timescales present in these young stars over this range of spectral types, giving rise to different braking efficiencies and histories.

1012.1817

1012

1012.3353_arXiv.txt

The cosmology of general fourth order corrections to Einstein gravity is considered, both for a homogeneous and isotropic background and for general tensor perturbations. It is explicitly shown how the standard cosmological history can be (approximately) reproduced and under what condition the evolution of the tensor modes remain (approximately) unchanged. Requiring that the deviations from General Relativity are small during inflation sharpens the current constraints on such corrections terms by some {\it thirty orders of magnitude}. Taking a more conservative approach and requiring only that cosmology be approximately that of GR during Big Bang Nucleosynthesis, the constraints are improved by $4\-- 6$ orders of magnitude.

\label{sec:1} General Relativity (GR) is often considered to be an extremely well tested physical theory, however compared to the other fundamental constants (e.g. the electron charge, the speed of light etc.), the gravitational constants are rather poorly constrained. Whilst it is true that the predictions of GR have been verified to staggering accuracy (see for example~\cite{C_Will}), there is little experimental/ observational evidence to restrict corrections of GR that involve higher powers of curvature invariants. This is simply a consequence of the fact that GR has been directly tested in weak curvature regimes, with strong curvature tests (such as near black holes and neutrons stars) typically resulting in theoretical restrictions. The Einstein-Hilbert action for GR contains only terms that are, at most, second order in the derivatives of the metric (in the form of the Ricci scalar $R$) and a natural question to then ask is what restrictions are there on the presence of terms that contain higher order derivatives of the metric? Motivation for such a question can be found from two, broadly different, approaches. The first is phenomenological: are there corrections to the Einstein-Hilbert action that can better describe observed data? In particular there has been much effort in looking for deviations from GR that might explain the apparent presence of Dark Energy, Dark Matter and inflation (such as $f(R)$ gravity~\cite{Nojiri:2006ri,Appleby:2007vb,Sotiriou:2008rp,Nojiri:2003ft,Amendola:2006we, Nojiri:2008nt,Nojiri:2010wj,Capozziello:2007ec,Capozziello:2006uv,Nojiri:2007as,Boehmer:2007kx}, chameleon models~\cite{Khoury:2003aq,Khoury:2003rn,Brax:2008hh}, conformal gravity~\cite{Mannheim:1999nc,Mannheim:1996rv}, MOND~\cite{Milgrom:1992hr,Milgrom:2002tu,Bekenstein:2007iq}, TeVeS~\cite{Skordis:2009bf,Moffat:2005si,Sagi:2007hb,Skordis:2005eu,Sanders:2005vd,Zlosnik:2006sb} etc.). The second approach is to consider theoretically motivated corrections to GR, often due to Quantum Gravity (such as String/ Brane theory~\cite{Mavromatos:2007sp,Smolin:1995ai,Baumann:2009ni,Brax:2003fv,Duff:2000mt,Garriga:1999yh, Nelson:2008sv,Germani:2002pt,Hawking:2000kj,Kachru:2003sx,Randall:1999vf,Sasaki:1999mi}, ADS-CFT~\cite{Anchordoqui:2000du,Gubser:1999vj,Nojiri:2000kq}, Loop Quantum Gravity~\cite{Bojowald:2006ww,lving_rev_MB,Smolin:2010kk, Nelson:2007cp, Bianchi:2009ri,Bianchi:2006uf,Smolin:2004sx,Ashtekar:2004eh,Thiemann:2007zz,Baez:1999sr}, Chern-Simons theory~\cite{Alexander:2009tp,Yunes:2008ua} etc.) or a Grand Unification theory (such as Non-commutative Geometry~\cite{Buck:2010sv,Chamseddine:2010ud,Chamseddine:2006ep}). Because there are only three curvature invariants that contain at most two derivatives of the metric, $R$, the cosmological constant $\Lambda$ and the Gauss-Bonnet combination (defined below), one is generally forced to consider correction terms containing $4^{\rm th}$ and higher derivatives of the metric\footnote{If these terms are built out of curvature invariants there is always an even number of derivatives.}, at least in the low energy, effective limit of the full theory. Theoretically there is, however, a serious difficultly with a theory that contains higher than second order derivatives of the metric; such a theory would contain `ghosts'. These are modes that have a negative kinetic term and can result in super-luminal propagation and other instabilities~\cite{Barth:1983hb,Stelle:1977ry, DeFelice:2006pg}. Such difficulties become particularly acute when the theory is quantised, since the presence of ghosts leads to particles with negative energies or states with negative norm. This shows that any theory containing higher order corrections to GR cannot be a fundamental theory unless powerful non-perturbative effects come into play. Here we will consider a general $4^{\rm th}$ order correction to GR and take the point of view that this is an {\it effective} theory, approximating (some of) the corrections to GR that are produced by some complete underlying, non-perturbative theory. This precisely is the approach that is taken when renormalisation group techniques are applied to GR~\cite{Lauscher:2001ya,Reuter:2001ag,Donoghue:1993eb}. The only terms containing at most $4^{\rm th}$ order derivatives of the metric that are constructed out of curvature invariants (other than the cosmological constant) are \be\label{eq:curv_inv} R^{\mu\nu\rho\gamma}R_{\mu\nu\rho\gamma}~,~~ R^{\mu\nu}R_{\mu\nu}~,~~ R^2~,~~ R~, \ee where $R^{\mu}_{~\nu\rho\gamma}$ is the Riemann tensor, defined as \be R^{\mu}_{~\nu\rho\beta} = \Gamma^{\mu}_{\rho\nu,\beta} - \Gamma^{\mu}_{\beta\nu,\rho} + \Gamma^{\alpha}_{\nu\rho} \Gamma^{\mu}_{\beta\alpha} - \Gamma^{\alpha}_{\nu\beta} \Gamma^{\mu}_{\rho\alpha}~, \ee and we follow the notation of~\cite{Nelson:2010ig} Greek letters to denote space-time indices; $\Gamma^{\mu}_{\nu\rho}$ is the usual Christoffel symbol; we denote partial derivatives as $\partial_\eta R = R_{,0}$ and we use the convention $R_{\mu\nu} = R^\rho_{~\mu\rho\nu}$, with the signature $\left( - + + +\right)$. The Ricci scalar (last term in (\ref{eq:curv_inv})) is at most second order in derivatives of the metric, while the particular combination, $R^2 - 4R_{\mu\nu}R^{\mu\nu}+R^{\mu\nu\rho\gamma}R_{\mu\nu\rho\gamma}$, called the Gauss-Bonnet combination, satisfies (in four dimensions) \be\label{eq:BG} \sqrt{-g}\left( R^2 - 4R_{\mu\nu}R^{\mu\nu} + R^{\mu\nu\rho\gamma}R_{\mu\nu\rho\gamma} \right) = {\rm total~ divergence}~. \ee Thus, assuming there are no boundary terms (which can be a significant complication, particularly when attempting to quantise the theory~\cite{Barth:1983hb}), one can write any $4^{\rm th}$ derivative correction to the Lagrangian of GR as \be\label{eq:action} {\cal S} = -\int {\rm d}^4 x \sqrt{-g} \left[ \frac{\gamma}{\kappa^2} R - \hbar \beta R^2 + \hbar \alpha R_{\mu\nu}R^{\mu\nu}\right] +{\cal S}_{\rm matter}~, \ee where for simplicity we have neglected the cosmological constant. Note that higher order terms (e.g. $R^3$) can lead to $4^{\rm th}$ derivative terms in the equation of motion~\cite{Schmidt:2001ac}, however here we will restrict our attention to terms that are of similar order in the Lagrangian, assuming that Eq.~(\ref{eq:action}) is a low curvature expansion of some underlying theory. We define $\kappa^2 = 32\pi G$ where $G$ is Newton's constant and we use units in which $\left[c\right]=1$. $\alpha$, $\beta$ and $\gamma$ are dimensionless couplings and the presence of $\hbar$ is due only to the dimensions (i.e.\ the theory is entirely classical). In general the coefficients $\alpha$ and $\beta$ are arbitrary constants, however if these corrections are due to some underlying Quantum Gravity theory, then they would be expected to be of order one, so that the quantum gravity scale is of the order of the Planck scale. The value of $\alpha$ and $\beta$ set the scale at which significant corrections to the GR occur and in the following we will restrict these values via cosmological considerations. In all of the following we set $\left[ \hbar\right]=1$. By varying this action with respect to the metric $g_{\mu\nu}$ as usual, one finds the field equations~\cite{Stelle:1977ry}, \beq\label{eq:EoM} H^\mu_{~\nu} &\equiv& \left( \alpha - 2 \beta\right) R^{;\mu}_{~~;\nu} - \alpha R^{\mu~;\rho}_{~\nu~;\rho} - \left( \frac{\alpha}{2} - 2 \beta \right) g^\mu_{~\nu} R^{;\rho}_{~;\rho} + 2\alpha R^{\rho\lambda} R^\mu_{~\rho \nu\lambda} \nonumber \\ && - 2\beta R R^\mu_{~\nu} - \frac{1}{2} g^\mu_{~\nu} \left( \alpha R^{\rho \lambda}R_{\rho\lambda} - \beta R^2\right) + \gamma \kappa^{-2} G^\mu_{~\nu} = -\frac{1}{2} T^\mu_{~\nu}~, \eeq where $R_{;\mu}=\nabla_\mu R$ is the covariant derivative of $R$ and $T_{\mu\nu}$ is the energy momentum tensor given by varying ${\cal S}_{\rm matter}$ with respect to $g_{\mu\nu}$. Perturbations of this theory around a flat background, show that the theory contains, in addition to the massless graviton, a massive spin-$2$ field (of negative energy) which is a ghost field and a (positive energy) massive scalar field~\cite{Stelle:1977ry}. One particular example of a theory that predicts $4^{\rm th}$ order corrections to GR is Non-commutative Geometry~\cite{Chamseddine:2010ud}. Here the asymptotic expansion of this (non-commutative) geometric theory produces the entire standard model coupled to a gravitational action of the form~\cite{Chamseddine:2006ep}, \be {\cal S}_{\rm NCG} = - \int {\rm d}^4 x \sqrt{-g} \left( \frac{1}{16\pi G}R + \alpha_{\rm NCG}C^{\mu\nu\rho\gamma}C_{\mu\nu\rho\gamma}\right)~, \ee where $C_{\mu\nu\rho\gamma}$ is the Weyl tensor\footnote{It should be mentioned that this Non-commutative Geometry theory is formulated in a Euclidean signature and is entirely classical.}. The consequences of this modification for cosmology~\cite{Nelson:2009wr,Nelson:2008uy} and astrophysics~\cite{Nelson:2010rt,Nelson:2010ru} have been considered and in particular in~\cite{Nelson:2008uy} it was shown that background cosmology is unaffected by the presence of such a correction. Using (\ref{eq:BG}) to write action this in the form of (\ref{eq:action}) gives, \be {\cal S}_{\rm NCG} = - \int {\rm d}^4 x \sqrt{-g} \left[ \frac{1}{16\pi G} R +2\alpha_{\rm NCG}\left( R_{\mu\nu}R^{\mu\nu}-\frac{1}{3}R^2\right)\right]~. \ee In particular then corrections of this form satisfy $\alpha = 3\beta$ and in the following we will show that {\it only} this combination leaves the background dynamics unaffected. The question we will address is what restrictions on $\alpha$ and $\beta$ can be deduced from the cosmology given by (\ref{eq:action}) and how they compare to the current constraints. Existing restrictions on $\alpha$ and $\beta$ are in fact very mild, requiring only that these coefficients be less than ${\cal O}\left( 10^{72}\right)$. The restrictions come from solar system tests, notable the perihelion precession of Mercury~\cite{Stelle:1977ry} and gravitational wave production in binary systems~\cite{Nelson:2010rt}. There are also laboratory scale tests of the inverse square law, which restrict parameters of this form significantly more, $\alpha < {\cal O} \left( 10^{60}\right)$~\cite{Kapner:2006si}. However here we will focus on the cosmological implications of such corrections and so compare to the other large scale constraints. The reason that these restrictions are so weak is that corrections to General Relativity occur at higher orders in curvature and hence are highly suppressed in weak curvature systems. In order for the effects of such corrections to become more significant, one needs to consider a system that contains either strong curvature (such as black holes or the early universe) or long evolution times (so as to allow the small deviations from General Relativity to accumulate). Here we consider the latter case, by examining both the dynamics of the background cosmology and the evolution of tensor mode perturbations over cosmological time scales. Ideally one would want to consider the evolution of scalar mode perturbations, since these can be directly related to the wealth of observational data coming from large scale galaxy surveys, weak lensing maps and the Cosmic Microwave Background (CMB). By contrast, cosmological tensor mode perturbations, although generically predicted by inflationary models, have yet to be observed. Despite this, tensor mode perturbations have the advantage of being technically much simpler to calculate than scalar modes (evolving according to a single equation, rather than four coupled equations). In addition, since tensor perturbations do not couple to the matter content of the universe other than through the background evolution (at least for matter with vanishing anisotropic stress) they are a direct probe of the underlying gravity theory and are unaffected by any possible modifications to the matter action. In Section~\ref{sec:BG} we derive the dynamics of the background (i.e.\ homogeneous and isotropic) cosmology, discussing in particular the solutions during inflation and during the radiation and matter dominated eras. Section~\ref{sec:tensors} derives the general evolution equation for tensor mode perturbations for the action given in (\ref{eq:action}), while in Section~\ref{sec:tensors_BG} we consider the evolution of such perturbations against specific background cosmologies. Using these evolution equations we derive constraints on $\alpha$ and $\beta$ in Section~\ref{sec:constraints} and summarise and conclude in Section~\ref{sec:conc}.

\label{sec:conc} We have shown that it is possible to approximately reproduce the key epochs of standard cosmology within a generally modified $4^{\rm th}$ order gravity theory. More precisely we demonstrated that {\it only} a correction to Einstein-Hilbert gravity of the form $C^{\mu\nu\rho\gamma}C_{\mu\nu\rho\gamma}$ introduces no deviation from standard background cosmology, while any other correction can approximate the standard background evolution, provided it is initially `close' to such a correction (see Section~\ref{sec:BG} for a more precise statement). Furthermore by requiring that the deviations from GR be small at certain key eras (such as during inflation and Big Bang Nucleosynthesis) we can quantify how `close' general corrections need to be in order to be consistent. Even taking the conservative point of view that radiation domination began only after the electroweak scale, such constraints are comparable to the best current restrictions on the theory. Further requiring that inflation be sourced by a fluid with $\omega \approx -1$ allowed us to improve this constraint {\it by thirty orders of magnitude}! The appearance of such large constraints is not uncommon within cosmology. Consider for example the spatial curvature within non-inflationary cosmology. If the energy density associated with such curvature vanishes to one part in $10$ today, it would have to vanish to one part in $10^{60}$ close to the Grand Unification scale. Indeed this was one of the original motivations for inflation. Similarly requiring that the higher order curvature corrections be small today is a {\it much} weaker constraint than requiring they be small in the past. The earlier we can examine the consequences of such corrections (such as during the radiation era or inflation) the more the constraints can be improved. At the background level we are able only to constrain the parameter $\alpha-3\beta$ ($\alpha=3\beta$ corresponds to the $C^{\mu\nu\rho\gamma}C_{\mu\nu\rho\gamma}$ correction of Non-commutative Geometry) and not the individual coefficients. In order to break this degeneracy we consider the evolution of perturbations against the homogeneous and isotropic background. For comparison with observations, one would like to examine the scalar perturbations of such a theory, which can be directly related to the galaxy power spectrum, weak lensing surveys, and the CMB. However the equations governing the evolution of scalar perturbation are extremely involved in a general $4^{\rm th}$ order gravity theory, making their use difficult. Instead we focused on the evolution of tensor modes (cosmological gravitational waves), which are technically simpler to calculate. While they have yet to be observed, they have the additional advantage of not coupling directly to matter (in the absence of anisotropic stress). Thus they are independent of any corrections to the matter action and are direct probes of the underlying geometry of the problem. Without direct observational evidence of such tensor mode perturbations, we considered the types of gravitational theories for which the evolution is approximately that of GR. The fact that such theories will typically lead to instabilities of perturbations, suggests that deviations from GR will include a strong growth of the amplitude of the perturbations. Had this occurred the perturbations would have been observed and hence our constraint is equivalent to ensuring that such instabilities occur at scales that are not (yet) observable. We showed that the presence of $R^2$ corrections to GR does not affect the evolution of tensor mode perturbations {\it in a flat radiation dominated universe}, while if {\it only} $R^2$ corrections are present, there is also no effect on the evolution, during the inflation era (although one can expect an affect on their production). Requiring that corrections be small for a specific mode and a particular scale (within the radiation era) hence constrains the coefficient of $R^{\mu\nu}R_{\mu\nu}$ corrections and breaks the degeneracy present in the background equations. By considering modes that have contributed to the measured (scalar) power spectrum\footnote{To avoid confusion we emphasise again, the tensor modes have {\it not} been measured. We consider these modes only because their scalar counterparts have been observed and quantified and if significant enhancement of the tensor modes had occur they would have contributed to the observed power spectrum at these scales.} we can estimate the constraint on $R^{\mu\nu}R_{\mu\nu}$ corrections, finding it to be improved by factor of $ 10^{4} \-- 10^{6}$ over current best restrictions. Thus, if inflation is to have occurred in the usual manner, we require \be \alpha - 3\beta \ll {\cal O}\left( 10^{40}\right)~, ~~~~ \alpha \ll {\cal O} \left( 10^{65}\right)~, \ee which is a remarkable improvement on current constraints which give $\alpha,~\beta \ll {\cal O}\left(10^{70}\right)$. Finally it is important to note that the presence of ghost in theories with such $4^{\rm th}$ order corrections is well known~\cite{Stelle:1977ry} and clearly indicates that the theory cannot be a fundamental theory (and certainly cannot be quantised in any standard way). Thus one should consider such corrections as an effective theory, which is valid in some range of scales and replaced by some more complete theory outside that range. For example, such corrections occur in the asymptotic expansion of certain Non-commutative Geometry theories~\cite{Chamseddine:2010ud}, (in that case $\alpha =3\beta$) and beyond a specific scale the space-time is no longer even approximately commutative. Here we have shown that the presence of ghosts is not felt at the level of background cosmology, however it is crucial for the evolution of perturbations and the constraints produced here essentially ensure that the presence of ghosts would not effect observable scales. One expects that higher (i.e.\ $6^{\rm th}$ and higher~\cite{Gottlober:1989ww,Hwang:1991}) order corrections to~(\ref{eq:action}) would become significant some scale. The presence of ghosts indicate the scale at which these higher order terms would need to become significant and essentially mark the scale at which the ($4^{\rm th}$ order) effective theory begins to break down. Thus the constraints produced here can be viewed as determining the scale at which (\ref{eq:action}) is a valid effective theory. It is important to note however, that here we have assumed a perturbative expansion and it is possible that non-perturbative effects might be significant.

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1012.1160_arXiv.txt

{We characterize the radial migration of stars in the disk plane by calculating the diffusion coefficient and the diffusion time-scale for a bulge-disk $N$-body self-consistent system with a marginally-stable Toomre-Q parameter. We find that diffusion is not constant in time, but follows the evolution of the bar, and becomes maximum near the corotation region and in the external disk region, where asymmetric patterns develop. \bigskip \noindent {\it Keywords} Galaxies: kinematics and dynamics, Galaxies: stellar content, Galaxies: spiral, Galaxy: disk}

The majority of local galaxies are barred-spiral galaxies, such as our Milky Way~\cite{delgado,eskridge} where the bar is the strongest non-axisymmetric pattern in the disk. The bar is generally coupled to other non-axisymmetric patterns such as spirals and warps if the disk is sufficiently cold~\cite{revaz}. The bar has a time-dependent activity, with a pattern speed which typically decreases in isolated galaxies. However, the system can be cooled again by adding dissipative infall of gas, or by forming stars on low-velocity dispersion orbits, with the net effect of restoring the amplitude of spiral waves and the strength of the bar, or even destroying it. In this way bars (and spiral waves) can be seen as recurrent patterns which can be rebuilt during their long history until the present configuration at $z=0$~\cite{bournaud}. Under the action of these non-axisymmetric patterns, stars move in the disk which gradually becomes hotter. Velocity dispersion of disk stars rises with age, as confirmed by observations in the Solar neighborhood (see~\cite{binney00} and references therein) and in external galaxies~\cite{gerssen, shapiro}. The origin and the amount of disk heating are still open to debate. First attempts to explain such a heating process in the disk of galaxies tried to model empirically the observed increase of the stellar velocity dispersion with age in the solar neighborhood. Wielen~\cite{wielen} suggested a diffusion mechanism in velocity space, which gives rise to typical relaxation times for young disk stars of the order of the period of revolution and to a deviation of stellar positions of 1.5~kpc in $200$~Myr. The result was obtained without making detailed assumptions on the underlying local acceleration process responsible for the diffusion of stellar orbits. Global acceleration processes, such as the gravitational field of stationary density waves or of central bars with constant pattern speed, were ruled out since their contribution to the velocity dispersion of old stars is negligible and concentrated in particular resonance regions~\cite{wielen,binneytremaine}. Different local accelerating mechanisms have been investigated so far in isolated galaxies, such as the gravitational encounters between stars and giant molecular clouds~\cite{spitzer51, spitzer53,lacey}, secular heating produced by transient spiral arms~\cite{barbanis,carlberg,fuchs} or the combination of the two processes~\cite{binney88, jenkins}. The radial excursion predicted by Wielen~\cite{wielen} is not sufficient to explain the weakness of the correlation between age and metallicity in the Solar neighborhood (see, for example, \cite{edvardsson}). In order to explain both the large scatter in the age-metallicity relationship and the evidence that even old disk stars today have nearly circular orbits, Sellwood \& Binney~\cite{sellwood} have recently suggested a new mechanism based on the resonant scattering of stars under the effect of transient spiral waves. In this process, a star initially on a nearly circular orbit resonates with a rotating wave and changes its angular momentum. If the duration of the peak amplitude of the perturbing potential is less than the period of the `horseshoe' orbits, i.e.~orbits of particles trapped at the corotation radius of the spiral wave, the star can escape from the potential well without changing its eccentricity. The net effect of this scattering mechanism is that stars migrate radially without heating the disk. In other words, the overall distribution of angular momentum is preserved, except near the corotation region of the transient spiral wave, where stars can have large changes of their angular momenta. The mechanisms driving radial migration and heating are still hotly debated. Minchev et al. ~\cite{minchev10a, minchev10b} propose a mixing mechanism where resonances between the bar and the spiral arms can act much more efficiently than transient spiral structure, dramatically reducing the predicted mixing timescales. Most of the observational signatures of radial mixing reported in the literature~\cite{grenon72, castro, grenon99} point to stars coming from a region next to the bulge/bar intersection. Radial migration of the stars in the disk have been suggested by Haywood~\cite{haywood}, who estimated upper values for the migration rate from 1.5 to 3.7 kpc/Gyr, which agree with the values estimated in~\cite{lepine} for the radial wandering due to the scattering mechanism assumed by Sellwood and Binney~\cite{sellwood}. High resolution cosmological simulations~\cite{roskar} confirm that such scattering mechanism determines a significant migration in the stellar disk. However these simulations ignore the important effects of bars found by Minchev et al. ~\cite{minchev10a, minchev10b}. Radial migration of stars (and gas) could have important implications for the interpretation of key observational constraints, such as the age-metallicity relationship or the metallicity gradients, since old stars that formed at small galactocentric radii from enriched gas or young metal-poor stars at large radii are enabled to appear in a Solar-neighborhood sample. Due to the lack of detailed information on the process driving radial mixing, models of the Galactic chemical evolution have evaluated past history of the solar neighborhood and the formation and evolution of the abundance gradients assuming that radial mixing did not played an important role~\cite{vandenbergh, schmidt, pagel, chiappini97, chiappini01}. Recently, Schoenrich and Binney~\cite{schonrich08} explored the consequences of mass exchanges between annuli by taking into account the effect of resonant scattering of stars described before. This approach appears to be successful to replicate many properties of the thick disk in the Solar neighborhood without requiring any merger or tidal event~\cite{schonrich09}. However, again in this case, the strong mixing mechanisms driven by bar resonances were not taken into account, casting thus doubts on some of their conclusions. In order to include the effect of radial migration in chemical evolution models and to gain a global (chemical + kinematical) vision of the processes at play in the galactic disks, many dynamical aspects need to be further investigated and in particular the role of the bar, that is the main non-axisymmetric component in disk galaxies. This is what we start to do in the present paper. In particular, we consider a marginally-stable disk which develops a central bar and spiral arms with the aim of quantitatively estimating the time and length scales of star diffusion in the radial direction. The final goal is to investigate how these characteristic scales evolve in time and how they depend on the activity of the bar. We present in detail the methods used for calculating the diffusion coefficient and time-scale. We will apply these methods to models which include the halo component and disks with different grade of stability in a forthcoming paper. The paper is organised as follows. In section~2 we describe the simulation and the relevant parameters. In section~3 we solve the diffusion equation in axisymmetric systems, we define the diffusion coefficient, the diffusion time-scale and the diffusion length-scale, and we describe how we estimate these quantities from the simulation results. In section~4 we present our results and in section~5 we summarise our findings.

Modeling the migration of stars in marginally stable disks as a diffusion process in the radial direction is a powerful tool which allows us to estimate quantitatively two crucial parameters, the diffusion coefficient and the diffusion time-scale. It is important to note that such diffusion model is only valid for describing the system for times smaller than the diffusion time-scale, which is of the order of the rotation period, and for times larger than the chaotic time-scale, after which orbits are no more time-reversible. The calculations of both the diffusion coefficient and the diffusion time-scale give us a quantitative measure of the migration process in the disk. Another advantage of studying the diffusion of stars in real space, rather than in velocity space, is that it can be easily related to the evolution of chemical elements, thus representing an interesting tool to be implemented in chemical evolution codes. We have found that the diffusion time-scale is of the order of one rotation period and that the diffusion coefficient $D$ depends on the evolution history of the disk and on the radial position. Larger values $D$ are found near the corotation region, which evolves in time, and in the external region, where asymmetric patterns develop. Marginally stable disks, with $Q_T \sim 1$, have two different families of bar orbits with different values of angular momentum $L_z$ and energy $E$, which determine a large diffusion in the corotation region. The application of the method described here to models which include the halo component and have disks with different initial Toomre-parameter $Q_T$ will be presented in a forthcoming paper. \smallskip \noindent {\bf \small Acknowledgements} {\small The simulations were run on the REGOR cluster at Geneva Observatory. This work has been supported by the Swiss National Science Foundation.}

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