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1012

1012.2286_arXiv.txt

Since 2006 WASP-South has been scanning the Southern sky for transiting exoplanets. Combined with Geneva Observatory radial velocities we have so far found over 30 transiting exoplanets around relatively bright stars of magnitude 9--13. We present a status report for this ongoing survey.

By 2004 several groups were pursuing systematic surveys for transiting exoplanets, including HAT (Bakos et\,al.\ 2004), OGLE (Udalski et\,al.\ 2002), TrES (O'Donovan et\,al.\ 2006) and XO (McCullough et\,al.\ 2006), and the WASP consortium had begun operating a ``SuperWASP'' camera array on La Palma in the Northern hemisphere (Pollacco et\,al.\ 2006). In the South the OGLE project were finding planets around stars of magnitude 15--16, but there was a clear opportunity for WASP's wide-field, bright-sky strategy, aimed at stars of magnitude 9--13. Funding from consortium universities was secured in 2004, leading to the construction of WASP-South at the South African Astronomical Observatory, closely copying Pollacco's design for the La Palma SuperWASP. Routine robotic operations commenced in March 2006. WASP-South is based around an array of 8 cameras, each with a 200-mm f/1.8 Canon lens backed by a {\it 2ev\/} 2048$\times$2048 Peltier-cooled chip. Each camera covers 7.8$^{\circ}\times$7.8$^{\circ}$, with 14$^{''}$ pixels, so that the array covers 15$^{\circ}\times$30$^{\circ}$. A broad-band filter gives a 400--700 nm bandpass. \begin{table}[t] \caption{The WASP-South planets} \begin{tabular}{lllllll} \tableline Name\rule{0mm}{3.5mm} & Position & Star & $V$ mag & Period & Mass & Radius \\ & & & & (days) & (Jup) & (Jup) \\ [1mm] \tableline WASP-4 \rule{0mm}{3.5mm} & 2334--42 & G7V & 12.6 & 1.34 & 1.12 & 1.42 \\ % WASP-5 & 2357--41 & G4V & 12.3 & 1.63 & 1.64 & 1.17 \\ % WASP-6 & 2312--22 & G8 & 12.4 & 3.36 & 0.50 & 1.22 \\ % WASP-7 & 2044--39 & F5V & {\ 9.5}& 4.95& 0.96 & 0.92 \\ % WASP-8 & 2359--35 & G6 & {\ 9.9}& 8.16& 2.24 & 1.04 \\ % WASP-15 & 1355--32 & F5 & 10.9& 3.75& 0.54& 1.43 \\ % WASP-16 & 1418--20 & G3V & 11.3& 3.12& 0.86& 1.01 \\ % WASP-17 & 1559--28 & F6 & 11.6& 3.74 & 0.49 & 1.74 \\ % WASP-18 & 0137--45 & F9 & {\ 9.3}& 0.94& 10.4& 1.17 \\ % WASP-19 & 0953--45 & G8V & 12.3& 0.79& 1.15& 1.31 \\ % WASP-20 & & & 10.7 & 2.4 & \\ WASP-22 & 0331--23 & & 12.0 & 3.53 & 0.56& 1.12 \\ % WASP-23 & 0644--42 & K1V & 12.7 & 2.94 & 0.87 & 0.96 \\ WASP-25 & 1301--27 & G4 & 11.9 & 3.76 & 0.58 & 1.22 \\ % WASP-26 & 0018--15 & G0 & 11.3 & 2.76 & 1.02 & 1.32 \\ % WASP-28 & 2334--01 & F9 & 12.0 & 3.41 & 0.91 & 1.12 \\ % WASP-29 & 2351--39 & K4V & 11.3 & 3.92 & 0.24 &0.79 \\ % WASP-30 & 2353--10 & F8V & 11.9 & 4.16 & 61.0 & 0.89 \\ % WASP-31 & 1117--19 & F7 & 11.7 & 3.41 & 0.47 & 1.56 \\ WASP-32 & 0015+01 & & 11.3 & 2.72 & 3.60 & 1.18 \\ % WASP-34 & 1101--23 & G5 & 10.4 & 4.32 & 0.59 & 1.22 \\ WASP-35 & & & 11 & 3.2 \\ WASP-36 & & & 12 & 1.5 & \\ WASP-37 & 1447+01 & G2 & 12.7 & 3.58 & 1.70 & 1.14 \\ % WASP-40 & & & 12 & 3 & \\ WASP-41 & 1242--30 & G8V & 11.6 & 3.05 & 0.93 & 1.21 \\ \tableline \end{tabular} \end{table} For the first two years WASP-South covered a ``zenith strip'' centred at --32$^{\circ}$, but skipping fields in the galactic plane that are too crowded. We typically raster 8--10 fields at a time with a cadence of $\sim$\,10 mins, taking two 30-sec exposures at each pointing. For the next two years WASP-South covered an equatorial strip centred at \mbox{--8$^{\circ}$}, partially overlapping with the Northern SuperWASP, while we are currently covering the far South with fields centred at --56$^{\circ}$. The data (approx 40 GB per night) are transferred on magnetic tape to Keele to be processed with the WASP pipeline. The data are then accumulated in an archive at Leicester where they are searched for candidate transits (see Collier-Cameron et\,al.\ 2006; 2007). The main limitations of the WASP design are the large pixels, often leading to blending, and the ``red noise'', which usually exceeds photon noise at magnitudes brighter than 12. Red noise is introduced by changes in focus, and by drift of the images over the chips through the night, by colour effects owing to the wide bandpass, and by variations in sky brightness such as through moonlight. For this reason the lightcurves are de-trended prior to the transit search.

1012.2286

1012

1012.5373_arXiv.txt

{} {The goal of our survey is to provide accurate and multi-epoch radial velocities, atmospheric parameters ($T_{\rm eff}$, $\log g$ and [M/H]), distances, and space velocities of faint red clump stars.} {We recorded high signal-to-noise (S/N$\geq$200) spectra of RC stars over the 4750-5950~\AA\ range at a resolving power 5500. The target stars are distributed across the great circle of the celestial equator. Radial velocities were obtained via cross-correlation with IAU radial velocity standards. Atmospheric parameters were derived via $\chi^2$ fit to a synthetic spectral library. A large number of RC stars from other surveys were re-observed to check the consistency of our results and the absence of offsets and trends.} {A total of 245 RC stars were observed (60 of them with a second epoch observation separated in time by about three months), and the results are presented in an output catalog. None of them is already present in other surveys of RC stars. In addition to astrometric and photometric support data from external sources, the catalog provides radial velocities (accuracy $\sigma$(RV)=1.3 km s$^{-1}$), atmospheric parameters ($\sigma(T_{\rm eff})$=88 K, $\sigma(\log g)$=0.38 dex and $\sigma$([M/H])=0.17 dex), spectro-photometric distances, (X,Y,Z) galacto-centric positions and (U,V,W) space velocities.} {}

The Red Clump (RC) is composed by low mass stars in the stage of central helium burning, following He ignition in an electron-degenerate core (Girardi 1999). They display properties that make them a primary tool to investigate Galactic structure and kinematics: \textit{(i)} their absolute magnitude shows minimal dispersion at optical and infrared wavelengths, \textit{(ii)} they are intrinsically bright, and thus observable throughout most of the Galaxy, \textit{(iii)} in magnitude-limited surveys they count for a fairly large fraction of observed targets, \textit{(iv)} their spectral types, ranging mainly from G8III to K2III, make them ideal stars to measure accurate radial velocities and atmospheric chemical abundances. Examples of recent applications of RC stars to Galaxy investigations are, among countless more, the peculiarities of Galactic rotation (Rybka et al. 2008), the stellar bar in the inner Galaxy (Cabrera-Lavers et al. 2007), the Galactic Bulge (Nataf et al. 2010), the vertical distribution of disk stars in terms of kinematic and metallicity (Soubiran et al. 2003), the surface mass density in the Galactic plane (Siebert et al. 2003), the origin of the Thick disk (Ruchti et al. 2010), the surface mass density in the Galactic disk (Bienayme et al. 2006) and age-metallicity relation (AMR), age-velocity relation (AVR) (Soubiran et al. 2008), tidal streams in solar neighborhood (Famaey et al. 2005, Antoja et al. 2008), Galactic substructures (Correnti et al. 2010, Law et al. 2010). The large proportion of RC stars observed by the ongoing RAVE survey (Steinmetz et al. 2006, Zwitter et al. 2008) and the accurate distances derived for them (Zwitter et al. 2010) support a great potential of the RAVE data base in progressing towards a better understanding of how the Galaxy formed, structured and evolved (Freeman \& Bland-Hawthorn 2002, Siebert et al. 2008, Veltz et al. 2008, Kiss et al. 2010). This paper is an extension to fainter magnitudes of the spectroscopic survey of RC stars of Valentini \& Munari (2010), hereafter named Paper I. Accurate radial velocities, atmospheric parameters ($T_{\rm eff}$, $\log g$, [M/H]), distances and space velocities have been obtained for 245 RC stars distributed along the great circle of celestial equator, 60 of them re-observed at a second epoch. These data are presented in a catalog together with photometric and astrometric support information from external sources. The application to Galaxy investigations of the results obtained in this study and in Paper I will be the topic of a forthcoming paper.

1012.5373

1012

1012.3719_arXiv.txt

We present a method for subtracting point sources from interferometric radio images via forward modeling of the instrument response and involving an algebraic nonlinear minimization. The method is applied to simulated maps of the Murchison Wide-field Array but is generally useful in cases where only image data are available. After source subtraction, the residual maps have no statistical difference to the expected thermal noise distribution at all angular scales, indicating high effectiveness in the subtraction. Simulations indicate that the errors in recovering the source parameters decrease with increasing signal-to-noise ratio, which is consistent with the theoretical measurement errors. In applying the technique to simulated snapshot observations with the Murchison Wide-field Array, we found that all 101 sources present in the simulation were recovered with an average position error of 10~arcsec and an average flux density error of 0.15\%. This led to a dynamic range increase of approximately 3 orders of magnitude. Since all the sources were deconvolved jointly, the subtraction was not limited by source sidelobes but by thermal noise. This technique is a promising deconvolution method for upcoming radio arrays with a huge number of elements, and a candidate for the difficult task of subtracting foreground sources from observations of the 21~cm neutral Hydrogen signal from the epoch of reionization.

\label{intro:Sect} The deconvolution of radio point sources is a problem that has been studied for several decades in radio astronomy. When calibration errors can be neglected, the problem of subtracting point sources from deconvolved radio images ultimately reduces to a problem of fitting their positions and flux densities as accurately as the instrumental noise permits. The methods used to deconvolve point source sidelobes are typically based on the CLEAN algorithm (Hogbom 1974; Clark 1980). The CLEAN algorithm looks for the brigthtest pixel in the image and subtracts a fraction of the dirty beam from the image at that location, forming a residual image. The search and subtraction loop is repeated until the sidelobes are reduced below the thermal noise level. The model components that are found through this iterative process can be convolved with a two dimensional Gaussian and introduced back into the residual image. The best estimate of flux density and position for each source is then found by fitting a two dimensional Gaussian to the source. The subtraction of point sources performed in this way has the known problem that the dynamic range achievable is limited by pixelization effects, i.e. by the fact that data are averaged and arranged into a regular grid. Therefore even a simple point source that does not lie at the centre of the grid cell cannot be represented by a single delta function model, but requires a potentially infinite number of components to be fully represented (Briggs \& Cornwell 1992, Perley 1999). In presence of the visibility data, the pixelization problem can be minimized and the dynamic range improved by subtraction of sources from the ungridded visibilities (Noordam \& de Bruyn 1982; Voronkov \& Wieringa 2004) and by centering the local pixel grid on the source to be deconvolved (Cotton \& Uson 2008). When the number of antenna elements to be correlated becomes extremely large, however, it becomes harder and harder to store the visibility data and the deconvolution has to be performed on images with, again, a limitation of the dynamic range due to pixelization effects. This is a relevant issue for upcoming radio telescopes like the Murchison Wide-field Array (MWA, Lonsdale et al. 2009) or future instrumentation like the SKA\footnote{http://www.skatelescope.org/} since they will produce a huge number of correlated visibilities. MWA will generate data at such a rate (approximately a PByte per day) that will be impractical to store the raw visibilities and go through the traditional selfcalibration loop, and the deconvolution of radio sources will happen in the image plane. The deconvolution of bright point sources is also a prominent issue in the view of the detection of the epoch of reionization (EoR) through the redshifted 21~cm line emission, which is one of the main goals of the MWA. The problem of foreground subtraction for EoR experiments has been studied by various authors in the literature (Di Matteo, Ciardi \& Miniati 2004; Morales \& Hewitt 2004; Santos, Cooray \& Knox 2005; Morales, Bowman \& Hewitt 2006; Wang et al. 2006, McQuinn et al. 2006; Gleser et al. 2008; Jeli{\'c} et al. 2008; Bowman, Morales \& Hewitt 2009; Liu et al. 2009a; Harker et al. 2009; Liu et al 2009b; Harker et al. 2010). Most of their efforts have been devoted to demonstrations that the diffuse Galactic synchrotron radiation and the classical confusion noise due to unresolved radio sources can be subtracted if it is assumed that they are spectrally smooth and absent of calibration errors. Recent observations (Ali, Bharadwaj \& Chengalur 2008; Bernardi et al. 2009; Pen et al. 2009; Parsons et al. 2010; Bernardi et al. 2010) have started to characterize the diffuse foreground component. All the simulations conducted so far, however, have assumed that the brightest point sources were perfectly subtracted from the data. Bowman et al. (2009) and Liu et al. (2009b) indicated that point sources should be subtracted down to a 10-100~mJy threshold in order to detect the EoR. Datta, Bhatnagar \& Carilli (2009) and Datta, Bowman \& Carilli (2010) studied the problem of subtraction of bright sources in the presence of calibration errors and concluded that sources brighter than 1~Jy should be subtracted with a positional precision better than 0.1~arcsec if calibration errors remain correlated over $\sim$6~hours of observation. If the errors are correlated on a shorter time length, however, they will tend to average down with time and the requirement for positional accuracy will be less stringent. Pindor et al. (2010) developed a technique based on matched filters to subtract bright point sources in MWA images in presence of diffuse emission. They showed that the dynamic range of the residual images can be improved by a factor of $\sim$2-3 in this way. In this paper we present a method of subtracting point sources from MWA dirty images that involves forward modeling and a nonlinear minimization scheme. Forward modeling is a general concept that can be used to extract astrophysical parameters from the data. We applied our method to simulated MWA images to show that point sources can be deconvolved with an accuracy limited by thermal noise even without storing the visibility data. The paper is organized as follows: in Section~\ref{forward_modeling} we present the method, in Section~\ref{various_applications} we apply the method to MWA simulated images and we conclude in Section~\ref{final_conclusions}

\label{final_conclusions} We have presented a point source deconvolution technique that makes use of forward modeling and an algebraic nonlinear minimization scheme. The main motivation for this implementation was achieving high dynamic range images in the absence of visibility data. Current (MWA) and future (SKA) radio interferometers require such a huge number of elements that they are being forced to rely more and more on real-time calibration and imaging, without the use of traditional selfcalibration techniques. The basic idea of our scheme is to forward model the sky brightness, i.e., to filter the sky model through the same instrumental response that is applied to the data. In the case of radio point sources, the forward model is the synthesized beam which is generated for each source individually. In this way, position dependent variations of the synthesized beam are accounted for. Point source astrophysical parameters are recovered through a nonlinear minimization over the image pixels. In this way we overcome the known dynamic range limitations of image-based deconvolution due to pixelization effects. Since the presented technique minimizes all the sources simultaneously and in an iterative way, it is minimally sensitive to sidelobe noise and essentially limited by thermal noise. It is worth noticing that this method can be applied to different sky components and can incorporate calibration parameters such as ionospheric displacements and primary beam shapes measured from the actual data. The technique was applied to three different simulated cases: a 10~min integration with the 32T MWA, an 8~sec snapshot image of the 512T MWA, and a 10~min integration with the 512T MWA where sources were placed inside and outside the field of view. In all cases we were able to subtract sources down to the thermal noise without assuming an a priori knowledge of the sky, with the exception of initializing the position and flux density of sources placed outside the field of view. The final residual images are consistent with the expected thermal noise on all the angular scales. Errors in the fitted parameters decrease with increasing SNR, in agreement with the expected theoretical measurement error distribution. Even when sources were not physically present in the images, we could subtract their sidelobes down to the thermal noise level. The 512T simulations are relevant in the light of the MWA EoR experiment. Since only a limited number of sources can be subtracted in real time, an off-line subtraction of the residual sources will have to be performed on the images to a high level of accuracy in order to precisely remove them and their direction-dependent synthesized beams. In the simulation of an 8~sec image with the 512T array, we achieved a dynamic range of $\sim$3400, indicating that the subtraction of foreground sources can be improved by 3 orders of magnitude through this technique. Source parameters can be retrieved with an average error of 10~arcsec on positions and 0.15\% errors on flux densities. The relative dynamic range of our subtraction is limited by the thermal noise and is above the 100~mJy threshold for 92\% of the sources. Since the best fit parameters improve with the SNR, a lower threshold - i.e. 10~mJy - can be reached by lowering the thermal noise through a longer integration. In fact, in the 512T 10~min simulation all the five sources present in the image had a dynamic range above the 10~mJy threshold, indicating that bright sources can be subtracted to a level that should not affect the detection of the EoR. A sky model more realistic than only point sources could be forward modeled by modifying the procedure presented here. Extended sky emission modeled as a list of delta functions (i.e. the equivalent of CLEAN components) could be directly treated by the present approach. More sofisticated modeling of extended emission that uses a set of basis functions like, for instance, shapelets (Yatawatta 2010) or a principal component analysis (de Oliveira-Costa et al. 2008) can be incorporated by convolving the model of the brightness distribution with the instrumental primary beam and then sampling it according to the $uv$ distribution (see Wayth et al. 2010 for an example of this approach). Future work will investigate these extensions and include a more realistic instrument model to better simulate the strategies for the EoR detection.

1012.3719

1012

1012.1146_arXiv.txt

High-resolution observations by visible and near-infrared interferometers of both single stars and binaries have made significant contributions to the foundations that underpin many aspects of our knowledge of stellar structure and evolution for cool stars. The CS16 splinter on this topic reviewed contributions of optical interferometry to date, examined highlights of current research, and identified areas for contributions with new observational constraints in the near future.

Observations of the fundamental parameters of stars - their masses, radii and effective temperatures - are crucial components in understanding stellar formation and evolution. For many types of stars, obtaining measurements with sufficient precision to constrain theoretical models can be challenging. A discussion of the contributions of optical interferometry to this field is timely given the technological maturation and recent scientific contributions of modern facilities such as VLTI, the Keck Interferometer, and the CHARA Array. Furthermore, a promising road lies ahead, as the 2-telescope measurements used in most of the work described below, are now being expanded by multi-way multi-channel combiners at the VLTI and CHARA which provide significant additional constraints by adding closure phase, differential phase, and in a few cases, imaging. Areas where optical interferometry is making compelling contributions towards parameterizing the fundamental properties of cool stars include: {\bf Radius} - The sub-milliarcsecond resolutions of modern optical interferometers allow for direct measurement of the linear sizes of nearby cool stars of all luminosity classes. {\bf Effective temperature} - Angular sizes, in combination with measured spectral energy distributions, provide direct quantification of this macroscopic quantity. {\bf Mass} - Dynamical masses from orbits determined with optical interferometers and radial velocity measurements are the highest-precision determinations of this most important fundamental parameter. Additionally, asteroseismological measurements of single stars provide direct measures of stellar mean density; coupled with interferometric radii, mass determinations are possible. {\bf Distance} - Parallactic measurements of stellar distances from optical interferometers calibrate stellar luminosities. {\bf Temperature Structure} - Limb-darkening measurements from optical interferometers probe the vertical temperature structure of a stellar atmosphere, providing constraints on model atmosphere structures in general, and models for convective flux transport in particular. Specific scientific questions that were considered during the splinter session were: (1) What have been the constraints on fundamental parameters for cool stars provided by optical interferometry? What are the limits on those constraints? (2) Where are the most attractive areas for guiding development of cool stars astrophysical models with observations from optical interferometers? (3) Are there specific areas of cool star evolution that are particularly well suited for studies by optical interferometry? (4) What are the observing opportunities available today for cool star research? (5) What are the prospects for future developments in optical interferometry that can significantly advance our knowledge of cool stars? Answers to these questions are of considerable interest to the Cool Stars 16 audience. Furthermore, the most recent event formally addressing fundamental stellar parameters was over a decade ago\footnote{IAU 189 in Sydney, Australia, ``Fundamental Stellar Properties: the Interaction Between Observation and Theory'' \citep{Bedding1997IAUS..189.....B}}, predating the current generation of operational facilities, and did not specifically address cool stars.

Optical interferometry is a challenging technique that rewards those who take up that challenge with observational data unobtainable via any other approach. In particular, contributions of the technique in the area of fundamental parameters of cool stars are substantial, highlighting areas of concern and interest for theorists. Furthermore, the facilities that provide access to this technique have matured substantially over the last 10 years, making it more accessible to all manners of astronomer.

1012.1146

1012

1012.0376_arXiv.txt

We present optical photometry of 16 transits of the super-Earth GJ~1214b, allowing us to refine the system parameters and search for additional planets via transit timing. Starspot-crossing events are detected in two light curves, and the star is found to be variable by a few percent. Hence, in our analysis, special attention is given to systematic errors that result from star spots. The planet-to-star radius ratio is $0.11610\pm 0.00048$, subject to a possible upward bias by a few percent due to the unknown spot coverage. Even assuming this bias to be negligible, the mean density of planet can be either $3.03\pm 0.50$~g~cm$^{-3}$ or $1.89\pm 0.33$~g~cm$^{-3}$, depending on whether the stellar radius is estimated from evolutionary models, or from an {\new{empirical mass-luminosity relation combined with}} the light curve parameters. One possible resolution is that the orbit is eccentric ($e \approx 0.14$), which would favor the higher density, and hence a much thinner atmosphere for the planet. The transit times were found to be periodic within about 15~s, ruling out the existence of any other super-Earths with periods within a factor-of-two of the known planet.

The recently discovered planet GJ~1214b (Charbonneau et al.~2009) is the smallest known exoplanet for which the mass, radius, and atmospheric properties are all possible to study with current technology. It is therefore a keystone object in the theory of planetary interiors and atmospheres, and has been welcomed as the harbinger of the ``era of super-Earths'' (Rogers \& Seager~2010). The planet's discoverers estimated the mass and radius of GJ~1214b to be $6.55\pm 0.98$~$M_\oplus$ and $2.68\pm 0.13$~$R_\oplus$, giving a mean density of $1.87\pm 0.40$~g~cm$^{-3}$ (Charbonneau et al.~2009). This is such a low density that it would seem impossible for the planet to be solid with only a thin, terrestrial-like atmosphere. Rather, it seems necessary to have a thick gaseous atmosphere, probably composed of hydrogen and helium but possibly also of carbon dioxide or water (Charbonneau et al.~2009, Rogers \& Seager 2010, Miller-Ricci \& Fortney 2010). In this paper, we report on observations of additional transits of GJ~1214b. As in other papers in the Transit Light Curve (TLC) series (Holman et al.~2006, Winn et al.~2007), one of our goals was to refine the basic system parameters, and thereby allow for more powerful discrimination among models of the planet's interior and atmosphere. Another goal was to check for any non-periodicity in the transit times, as a means of discovering other planets in the system, through the method of Holman \& Murray (2005) and Agol et al.~(2005). Super-earths have frequently been found in pairs or even triples in compact arrangements (Lo Curto et al.~2010), and it would be interesting to know if GJ~1214b is another such example. This paper is organized as follows. Section \ref{sec:obsred} describes the observations and data reduction. Section~\ref{sec:model} presents the light curve model, taking into account the effects of starspots. Section~\ref{sec:analysis} discusses the method by which we estimated the model parameters and their confidence intervals. Section~\ref{sec:radrat} discusses the results for the planet-to-star radius ratio. Section~\ref{sec:radplanet} presents two different methods for determining the stellar radius (and hence the planetary radius), which give discrepant results. Some possible resolutions of this discrepancy are discussed. Section~\ref{sec:timing} presents our analysis of the measured transit times, and constraints on the properties of a hypothetical second planet. Finally, in Section~\ref{sec:disc}, we discuss the implications of our analysis on the understanding of GJ~1214b and more broadly on M dwarf transit hosts.

\label{sec:disc} One of our goals in this study was to improve on the estimates of the basic parameters of GJ~1214b. Yet despite having undertaken many high-precision observations of transits, we have not significantly improved on the estimate of the planetary radius. In fact\ron{,} our analysis has led us to conclude that the radius is even more uncertain than had been recognized previously. This is for two reasons. First, the clear evidence for starspots in our most precise light curves has caused us to consider the possible effects of stellar activity on the analysis of transit photometry. Second, and more importantly, we have found a significant disagreement between two methods of estimating the stellar and planetary dimensions, with no good reason why either one should be disregarded or considered less trustworthy. Both of these complications were known prior to this work {\new{(e.g., Charbonneau et al. 2009)}}, but we have brought them into {\new{sharper}} focus. The problem of star spots can be mitigated by observing in longer-wavelength bandpasses. This is because the flux contrast between two blackbodies of different temperatures is a decreasing function of wavelength. {\new{For example, if we assume}} the particular starspots that produced anomalies in our data are approximately the same size as the planet, then they are $\approx$150~K cooler than the stellar photosphere (see \S~\ref{sec:spotcross} for details). This corresponds to a surface-brightness ratio of about 0.67 between the spots and the surrounding \ron{photosphere}, at an observing wavelength of 0.6~$\mu$m. At 3.5~$\mu$m, the surface-brightness ratio would be about 0.91, representing a smaller contrast and correspondingly smaller starspot-induced effects. The problem of the stellar radius will be more difficult to solve, and there is much at stake. The mean planetary density could be $1.89 \pm 0.33$~g~cm$^{-3}$ or $3.03 \pm 0.50$~g~cm$^{-3}$, depending on which route is taken to estimate the stellar radius. The lower density would imply that the planet must have a dense gaseous atmosphere, for which there are many intriguing possible origins and compositions (Rogers \& Seager 2010, Miller-Ricci \& Fortney 2010). In contrast, the higher density could be consistent with a solid planet with a very thin (or nonexistent) atmosphere. We have discussed several possible resolutions of this discrepancy, and argued that the most attractive possibility is that the planet has an eccentric orbit, $e \approx 0.14$. This hypothesis can be tested by gathering additional RV data, or by measuring the time of occultation of the planet by the star. Neither task will be easy. Assuming the minimum detectable eccentricity to vary inversely as the square root of the number of RV data points, one would need approximately 4 times as many data points (with the same precision as the existing data) for a 2$\sigma$ detection of an eccentricity of 0.14. Likewise, the occultation depth is expected to be smaller than 0.1\% at 3.5~$\mu$m, the region accessible to the {\it Spitzer} space telescope, which probably offers the best prospects for such a detection. This experience with GJ 1214 invites some broader remarks about the suitability of M dwarf stars as targets in surveys for transiting planets. The {\new{advantages of M dwarfs have been discussed}} by Nutzman \& Charbonneau (2009), among others. Their smaller radii and masses allow for smaller planets to be detected, at a given signal-to-noise ratio. They are numerous in our Galactic neighborhood. Transits of planets in the habitable zone in particular are more probable {\new{and, hence}}, more frequent than they are around higher-mass stars. These are decisive advantages, fully justifying the ongoing efforts that concentrate exclusively on M dwarfs. However, there are two disadvantages of low-mass stars. They tend to have larger spots and an overall higher level of stellar activity than more Sun-like stars, which will interfere with the measurement of the basic transit parameters (c.f.\ Seager \& Deming 2009). And, it will be harder to obtain reliable estimates of the stellar mass and radius, because of possible limitations of stellar evolutionary models in the range 0.3--0.7~$M_\odot$ and, more generally, because of our more limited empirical knowledge of the lowest-mass stars.

1012.0376

1012

1012.5079_arXiv.txt

Numerical aspects of dynamos in periodic domains are discussed. Modifications of the solutions by numerically motivated alterations of the equations are being reviewed using the examples of magnetic hyperdiffusion and artificial diffusion when advancing the magnetic field in its Euler potential representation. The importance of using integral kernel formulations in mean-field dynamo theory is emphasized in cases where the dynamo growth rate becomes comparable with the inverse turnover time. Finally, the significance of microscopic magnetic Prandtl number in controlling the conversion from kinetic to magnetic energy is highlighted.

There are two important aspects connected with astrophysical dynamos compared with dynamos on a bicycle. Firstly, they are self-excited and do not require any permanent magnets. Secondly, they are {\it homogeneous} in the sense that the medium is conducting everywhere in the dynamo proper and there are no wires or insulators inside. Self-excited dynamos were invented by the Danish inventor S{\o}ren Hjorth, who received the patent for this discovery in 1854, some 12 years before Samuel Alfred Varley, Ernst Werner von Siemens and Charles Wheatstone announced such an invention independently of each other. Von Siemens is known for having recognized its industrial importance in producing the most powerful generators at the time, for which he, in turn, received a patent in 1877. The idea that homogeneous dynamos might work in the Sun, was first proposed by \cite{Lar19} in a one-page paper. However, some 14 years later, \cite{Cow33} showed that axisymmetric dynamos cannot work in a body like the Sun. At the time it was not clear whether this failure was genuine, or whether it was critically connected with Cowling's assumption of axisymmetry. The suspicion that the third dimension might be critical was not particularly emphasized when \cite{Lar34} tried to defend his early suggestion with the words ``the self-exciting dynamo analogy is still, so far as I know, the only foundation on which a gaseous body such as the Sun could possess a magnetic field: so that if it is demolished there could be no explanation of the Sun's magnetic field even remotely in sight.'' The essential idea about the operation of the solar dynamo came from \cite{Par55}, who developed the notion that cyclonic events would tilt a toroidal field systematically in the poloidal direction, closing thereby a critical step in the dynamo cycle. While this concept is still valid today, it still required the existence proof by \cite{Her58} that began to convince critics that Cowling's antidynamo theorem does not extend to the general case of three dimensions. Nevertheless, subsequent progress in modeling the solar dynamo appears to have been suspended until the foundations of a mean-field treatment of the induction equation were developed by \cite{SKR66}. In the following years, a large number of models were computed covering mostly aspects of the solar dynamo \citep{SK69a,Par70,Par70b,Par70c,Par71b,Par71d,Par71f}, but in some cases also terrestrial dynamos \citep{SK69b,Par71c} and the galactic dynamo \citep{Par71,Par71e,VR71,VR72}. These developments provided a major boost to dynamo theory given that until then work on the galactic dynamo, for example, focussed on aspects concerning the small-scale magnetic field \citep{Par69}, but not the global large-scale fields on the scale of the entire galaxy. In fact, also regarding small-scale dynamos, there were important developments made by \cite{Kaz68}, but they remained mostly unnoticed in the West, even when the first direct simulations by \cite{MFP81} demonstrated the operation of such a dynamo in some detail. In fact, in some of these dynamos, the driving of the flow involved helicity, but its role in helping the dynamo remained unconvincing, because no large-scale field was produced. We now understand that this was mainly because there was not enough scale separation between the scale of the domain and the forcing scale, and that one needs at least a ratio of 3 \citep{HBD04}. Simulations in spherical geometry were much more readily able to demonstrate the production of large-scale magnetic fields \citep{Gil83,Gla85}, but even today these simulations produce magnetic fields that propagate toward the poles \citep{KKBMT10} and not toward the equator, as in the Sun. We can only speculate about possible shortcomings of efforts such as these that must ultimately be able to reproduce the solar cycle. Several important developments happened in the 1980s. Firstly, it became broadly accepted that the magnetic field inside the Sun might be in a fibril state \citep{Par82}, i.e.\ the filling factor is small and most of the field is concentrated into thin flux tubes, as manifested by the magnetic field appearance in the form of sunspots at the surface. However, such tubes would be magnetically buoyant, and are expected to rise to the surface on a time scale of some 50 days \citep{MI83,MI86}. This time is short compared with the cycle time and might lead to excessive magnetic flux losses, which then led to the proposal that the magnetic field would instead be generated in the overshoot layer beneath the convection zone. This idea is still the basic picture today, although simulations of convection generally produce magnetic fields that are distributed over the entire convection zone. Yet another important development in the 1980s was the proposal that the $\alpha$ effect might actually be the sum of a kinetic and a magnetic part and that the magnetic part can be estimated by solving an evolution equation for the magnetic helicity density. The importance of this development was obscured by the excitement that the two evolution equations for poloidal and toroidal field, supplemented by a third equation for the magnetic helicity density, could produce chaos \citep{Ruz81}. The connection to what was to come some 10--20 years later was not yet understood at that point. Simulations of helical MHD turbulence in a periodic domain demonstrated that in a periodic domain the $\alpha$ effect might be quenched in an $\Rm$-dependent fashion like \EQ \alpha={\alpha_0\over1+\Rm\meanBB^2/\Beq^2}. \EN If this were also true of astrophysical dynamos, $\alpha$ would be negligibly small and would not be relevant for explaining the magnetic field in these bodies. Such quenching is therefore nowadays referred to as catastrophic quenching. However, there is now mounting evidence that this type of $\alpha$ quenching is a special case of a more general formula \citep[see, e.g.,][]{B08} \EQ \alpha={\alpha_0+\Rm\left[ \eta_{\rm t}\mu_0\meanJJ\cdot\meanBB/B_{\rm eq}^2 -(\nab\cdot\meanFF_{\rm C})/2 k_{\rm f}^2B_{\rm eq}^2 -(\partial\alpha/\partial t)/2\eta_{\rm t} k_{\rm f}^2\right] \over1+\Rm\meanBB^2/B_{\rm eq}^2}, \label{QuenchExtra2} \EN which comes from magnetic helicity conservation. Note that in this equation there are 3 new terms that all scale with $\Rm$ and are therefore important. Even in a closed or periodic domain, the first and third terms in squared brackets contribute, the most promising way out of catastrophic quenching is through magnetic helicity fluxes \citep{BF00,Klee00}. These developments are still ongoing and we refer here to some recent papers by \cite{SSSB06,BCC09,Can10}.

Simulating simple dynamos on the computer is nowadays quite simple. Nevertheless, we have seen here examples that illustrate that things can also go quite ``wrong''. In the case of magnetic hyperdiffusion, it is clear what happens \citep{BS02}, so that magnetic hyperdiffusion can also be used to ones advantage, as was demonstrated in \cite{BDS02}. However, in the case of Euler potentials it is not clear what happens and whether this method can be used to simulate even the ideal MHD equations, given that each numerical scheme will introduce some type of diffusion. In this short review, we have also attempted to clarify why numerical calculations of $\alpha$ effect and turbulent diffusion using the standard test-field method \citep{Sch07} would yield values that can only reproduce a correct growth rate in the case of vanishing growth. In all other case, a representation in terms of integral kernels has to be used. Finally, we have discussed some effects of using magnetic Prandtl numbers that are different from unity. It turns out that in the steady state, the rate of transfer from kinetic to magnetic energy depends on the value of $\Pm$. This is somewhat unexpected, because the onset condition for dynamo action does not depend on $\Pm$ \citep{B09}, and yet the actual efficiency of the dynamo, as characterized by the work done against the Lorentz force, $-\bra{\UU\cdot(\JJ\times\BB)}$, does depend on $\Pm$ and is proportional to $\Pm^{-n}$ (with $n$ between 1/2 and 2/3) for large values of $\Pm$. Understanding the limits of numerical simulations is just as important as appreciating its powers. As the example with the problem with mean-field and simulated growth rates shows, understanding the initial mismatch can be the key to a more advanced and more accurate theory that will ultimately be needed when describing some of the yet unexplained properties of astrophysical dynamos.

1012.5079

1012

1012.2235_arXiv.txt

We present the ellipticity distribution and its evolution for early-type galaxies in clusters from $z\sim0.8$ to the current epoch, based on the WIde-field Nearby Galaxy-cluster Survey (WINGS) ($0.04 \leq z \leq 0.07$), and the ESO Distant Cluster Survey (EDisCS) ($0.4 \leq z \leq 0.8$). We first investigate a mass limited sample and we find that, above a fixed mass limit ($M_{\ast}\geq 10^{10.2}M_{\odot}$), the ellipticity ($\epsilon$) distribution of early-types noticeably evolves with redshift. In the local Universe there are proportionally more galaxies with higher ellipticity, hence flatter, than in distant clusters. This evolution is due partly to the change of the mass distribution and mainly to the change of the morphological mix with z (among the early types, the fraction of ellipticals goes from $\sim70\%$ at high-z to $\sim40\%$ at low-z). Analyzing separately the ellipticity distribution of the different morphological types, we find no evolution both for ellipticals and for S0s. However, for ellipticals a change with redshift in the median value of the distributions is detected. This is due to a larger population of very round ($\epsilon <0.05$) elliptical galaxies at low-z. In order to compare our finding to previous studies, we also assemble a magnitude-``delimited'' sample that consists of early-type galaxies on the red sequence with $-19.3 >M_{B} +1.208z >-21$. Analyzing this sample, we do not recover exactly the same results of the mass-limited sample.his indicates that the selection criteria are crucial to characterize the galaxy properties: the choice of the magnitude-``delimited'' sample implies the loss of many less massive galaxies and so it biases the final conclusions. Moreover, although we are adopting the same selection criteria, our results in the magnitude-``delimited'' sample are also not in agreement with those of \cite{h09}. This is due to the fact that our and their low-z samples have a different magnitude distribution because the \cite{h09} sample suffers from incompleteness at faint magnitudes.

Ellipticals and lenticulars (S0s) belong to the class of early-type galaxies. This means that they have several properties in common: they dominate the total galaxy population at high masses, they preferentially inhabit dense regions of the universe, such as rich clusters \citep{dressler97}, they tend to be passive, they have red colors and their spectra show strong values of the characteristic $D4000$ feature (see e.g. \citealt{kauffmann03, brinchmann04}); they lack spiral arms and in most cases exhibit neither major dust features nor a large interstellar gas content. For these reasons, often they are considered together. On the other hand, elliptical and S0 galaxies differ in several important ways: S0s are bulge-dominated systems with an identifiable disk (e.g. \citealt{scorza98,laurikainen07}), that is mainly rotationally supported (e.g. \citealt{erwin03, cappellari05}), their intrinsic shape is similar to that of spirals \citep{rood67, sandage70} and their formation is still not well understood. \cite{hubble36} first proposed their existence as transitional class between ellipticals and spirals. Understanding how they form and evolve is essential if we wish to have a complete picture of how galaxy morphology is related to galaxy formation and the environment. Then again, % ellipticals show ellipsoidal shapes, not rarely with significant kinematic twists, and kinematically decoupled components in their centres. Most of them are not characterized by strong rotation \citep{bertola75}, and their luminosity profiles follow a Sersic's law. In the Local Universe disky ellipticals are probably the high bulge mass end of S0. Morphologically, \cite{dressler97} showed that, at least for bright galaxies, the % raising fraction of early-type galaxies since $z\sim 0.5$ corresponds mainly to an increase of lenticular S0 galaxies, with a roughly constant elliptical fraction. S0s are quite rare in clusters at high redshift ($z>0.3-0.4$); as a consequence, they have to acquire their shapes with different time-scales and later than ellipticals. The evolving fraction of S0s in clusters might result from the evolving population of newly accreted spiral galaxies from infalling groups and the field. \cite{fasano00} showed that the cluster S0 to elliptical ratio is, on average, a factor of $\sim 5$ higher at $z\sim 0$ than at $z\sim0.5$. At higher redshift, there is no evidence for any further evolution of the S0 fraction in clusters to $z\sim 1$: most of the evolution occurs since $z \sim 0.4$ (see e.g. \citealt{postman05, desai07, wilman09}). \cite{dressler97} and \cite{postman05} also investigated the ellipticity distributions of the S0 and elliptical galaxies in their magnitude limited samples. They found that the ellipticity distribution of S0 and elliptical galaxies show no evolution over the broad redshift ranges in their samples. Moreover, they differ from each other, providing evidence for the existence of two distinct classes of galaxies. In contrast, in their magnitude-``delimited'' sample (with both an upper and a lower magnitude limit), \cite{h09} found no evolution in neither the median ellipticity nor the shape of the ellipticity distribution with redshift for early-type (ellipticals + S0s) red-sequence galaxies. This lead them to conclude that there has been little or no evolution in the overall distribution of bulge-to-disk ratio of early-type galaxies from $z\sim 1$ to $z\sim 0$. Assuming that the intrinsic ellipticity distribution of both elliptical and S0 galaxies separately remains constant, they finally concluded that the relative fractions of ellipticals and S0s do not evolve from $z\sim 1$ to $z = 0$ for a red-sequence selected sample of galaxies. All the cited works analyzed samples limited in some ways by magnitude cuts. For the first time, in this paper we analyze the evolution of the ellipticity distribution of early-type galaxies also in a mass-limited sample. For the sample in the Local Universe, we analyze the data of the WIde-field Nearby Galaxy-cluster Survey (WINGS) \citep{fasano06}, while for that in the distant Universe we use the ESO Distant Cluster Survey (EDisCS) \citep{white05}. These large cluster samples and their high quality images (see \S 2) allow us to characterize properly the cluster environment at the two redshifts and to subdivide galaxies into the different morphological types and obtain robust estimates of ellipticity. This paper is organized as follows: in \S 2 we present the cluster and galaxy samples (WINGS \citep{fasano06} and EDisCS \citep{white05}), describing the surveys, the data reduction, the determination of morphologies, ellipticites and masses. We also depict the selection criteria we follow to assemble the mass-limited and the magnitude-``delimited'' samples. In \S 3 we show the results of our analysis of the evolution of the ellipticity distribution with redshift in our mass-limited samples, while in \S 4 we show the same for the magnitude-``delimited'' samples. In \S 5 we try to reconcile the results of the different samples, while in \S 6 we compare our results with those found in literature (in particular with the results drawn by \citealt{h09}). Finally, in \S 7 we discuss and summarize our findings. Throughout this paper, we assume $H_{0}=70 \, \rm km \, s^{-1} \, Mpc^{-1}$, $\Omega_{m}=0.30$, $\Omega_{\Lambda} =0.70$. The adopted initial mass function is a \cite{kr01} in the mass range 0.1-100 $M_{\odot}$.

In this paper we have analyzed the ellipticity distribution of early-types galaxies, and of ellipticals and S0s separately, in clusters at $z=0.04-0.07$ and $z=0.4-0.8$. We have taken into account both a mass-limited sample and a magnitude-``delimited'' sample of galaxies. \begin{itemize} \item In our mass-limited samples, above the common mass limit ($M_{\ast}\geq 10^{10.2}M_{\odot}$) the ellipticity distribution of early-type galaxies strongly varies with redshift. This is due to a change both of the median and of the shape of the distributions with redshift. For ellipticals, no statistically significant differences are observed in the high- and low-z distribution, even if an evolution of the medians is detected and we observe an excess population of round ellipticals at low-z compared to high-z. Finally, no evolution is observed for S0s. The evolution of early-type galaxies is not simply related to the different mass distributions at high- and low-z. In fact, removing the influence of the mass, the results remain inconsistent. Instead, it is mainly related to the evolution of the morphological mix with redshift and hence to the relative contribution of ellipticals and S0s at the two epochs. \item As mentioned in the previous point, in our low-z sample, we find a population of very round ($\epsilon\leq0.05$) elliptical galaxies that is less conspicuous at high-z. This population seems real and not due to selection effects or measurement problems. \item In our magnitude-``delimited'' sample, for early-types and S0s the evolution is not evident (though the medians of both early types and S0s change with z), while for ellipticals we have found a change of the distribution with redshift. \item The observed differences between the mass-limited sample and the magnitude-``delimited'' one can be due to the different mass distribution of the two samples: in fact in the magnitude-``delimited'' samples we are loosing some galaxies that enter the mass-limited one, both at high but especially and more importantly at low masses. \item Our magnitude-''delimited'' results are not in agreement with those of \cite{h09}, who also analyzed a magnitude-``delimited'' sample of early-types belonging to the red sequence. The main reason of the observed discrepancy is that, despite galaxies being selected following in principle the same criteria, in practice the two low-z samples have a different magnitude distribution because the \cite{h09} sample suffers from incompleteness at faint magnitudes. \end{itemize}

1012.2235

1012

1012.4024_arXiv.txt

{ If the positron fraction and combined electron-positron flux excesses recently observed by PAMELA, Fermi and HESS have a dark matter origin, final state radiation (FSR) photons from dark matter annihilation into lepton-rich final states may be detected with observations of satellite dwarf galaxies of the Milky Way by ground-based atmospheric Cherenkov telescopes (ACTs). We find that current and near-future ACTs have excellent potential for such detection, although a discovery cannot be guaranteed due to large uncertainties in the distribution of dark matter within the dwarfs. We find that models predicting dark matter annihilation into two-lepton final states and those favoring four-lepton final states (as in, for example, ``axion portal" models) can be reliably distinguished using the FSR photon spectrum once measured, and the dark matter particle mass can also be accurately determined.} \FullConference{Identification of Dark Matter 2010\\ July 26 - 30 2010\\ University of Montpellier 2, Montpellier, France} \begin{document}

Recent measurements of the positron fraction in cosmic rays in the 10--80 GeV range by PAMELA~\cite{pamela} and the combined electron-positron flux up to the TeV scale by Fermi~\cite{fermi} and HESS ~\cite{hess} indicate excesses inconsistent with conventional astrophysical background. These excesses can be explained, among other possibilities, by dark matter with an annihilation cross section of $\langle \sigma v\rangle \sim3\times10^{-23}$ cm$^3$s$^{-1}$ in the Milky Way and a mass of 1--3 TeV, provided annihilation is predominantly into electrons or muons~\cite{patrick,bestfits}. If these excesses indeed have a dark matter origin, accompanying signals are expected in the form of energetic gamma rays. For leptophilic dark matter, the high energy end of the gamma ray signal is dominated by final state radiation (FSR). Dwarf galaxies --- made up almost entirely of dark matter, with no detected neutral or ionized gas, minimal dust, no magnetic fields, and little or no recent star formation activity --- are favorable targets for searches for such gamma rays. And since dark matter gamma ray signals from dwarf galaxies are mainly constrained by low statistics, atmospheric Cherenkov telescopes (ACTs), with typical effective areas $\sim 10^5$ times that of Fermi, offer a distinct advantage. In this work, we focus on the prospects of detecting FSR from dark matter annihilation with ACT observations of dwarf galaxies. We also investigate the important question of whether the ACTs can measure the FSR photon spectrum precisely enough to distinguish between different leptophilic dark matter models. For details, please refer to \cite{fsrpaper}. \textit{Final state radiation (FSR) ---} FSR is present whenever dark matter annihilates to charged particles, as is the case for the following three leptophilic models motivated by fits to PAMELA, Fermi, and HESS data \cite{patrick,bestfits} (with dark matter denoted by $\chi$):\\ (i) Model A: $\chi\chi\rightarrow\mu^+\mu^-$.~~~~ (ii) Model B: $\chi\chi\rightarrow\phi\phi\rightarrow4e$.~~~~(iii) Model C: $\chi\chi\rightarrow\phi\phi\rightarrow4\mu$.\\ $\phi$ denotes an intermediate ``portal" particle, with mass taken to be of order 1 GeV. Factorization theorems ensure that the energy spectrum of the FSR photons is, to leading order in $m_l/m_{\chi}$, independent of the details of the annihilation process, allowing for quasi-model-independent predictions. For annihilation into a lepton-antilepton pair $l\bar{l}$ as in model A, the FSR flux within the leading log approximation is \cite{robust} \beqa \frac{d\Phi_{FSR}}{dx}=\Phi_0\left(\frac{\left<\sigma v\right>}{1pb}\right)\left(\frac{100~{\rm GeV}}{m_{\chi}}\right)^3\frac{1+(1-x)^2}{x}\,\log\left(\frac{4m_\chi^2(1-x)}{m_l^2}\right)\,J, \label{2state}\\ J=\frac{1}{8.5~{\rm kpc}}\left(\frac{1}{0.3~{\rm GeV/cm}^3}\right)^2L,~~~~L= \int d\Omega\int_{l.o.s.}\rho^2 dl. \label{jfactor} \eeqa{factors} Here $x=2E_{\gamma}/\sqrt{s}=E_{\gamma}/m_{\chi}$, $\Phi_0=1.4\times 10^{-14}$ cm$^{-2}$s$^{-1}$GeV $^{-1}$, and $J$ is a dimensionless factor that carries all the astrophysics information;\footnote{Eq. \leqn{2state} and similar formulas below are equally applicable for decaying dark matter, with an appropriate redefinition of the $J$ factor. However, we do not consider decaying dark matter in this paper since the resulting FSR signals from dwarf galaxies are generally too weak to be detected.} astrophysical uncertainties, therefore, do not affect the spectrum of final state radiation. The spectrum in Eq. (\ref{2state}) features a characteristic ``edge" at the dark matter mass \cite{robust}. For 4-body annihilation as in models B and C, the FSR spectrum is \cite{robust} \beq \frac{d\Phi_{FSR}}{dx}=\Phi_0\left(\frac{\left<\sigma v\right>}{1pb}\right)\left(\frac{100~{\rm GeV}}{m_{\chi}}\right)^32\frac{2-x+2x\log x-x^2}{x}\,\log\left(\frac{m_\phi^2}{m_l^2}\right)\,J. \eeq{4state} For models A and C, annihilation is to muons, and final state radiation form the subsequent decay of the muon should also be taken into account. Relevant formulas for this contribution are as given in \cite{rouven, fsrpaper}. Typically, FSR off muons from the annihilation process remains dominant unless $m_{\phi}\sim m_{\mu} $\cite{rouven}. In this work we fix our parameters to $\left<\sigma v\right>=3\times10^{-23}$ cm$^3$s$^{-1}$, $m_{\chi}=3$ TeV, and $m_{\phi}=1$ GeV, as favored by fits to PAMELA, Fermi, and HESS data. It should be kept in mind that lower velocities in dwarf galaxies can lead to larger cross sections via greater Sommerfeld enhancement. \textit{Dwarf galaxies ---} We use the following dwarf galaxies, which have been known to be promising candidates for dark matter searches, in our analysis: \\ \indent Draco (18.63~$\pm$~0.60) ~~~~~~~~~ Ursa Minor (18.79~$\pm$~1.76)\\ \indent Willman 1 (19.55~$\pm$~0.98) ~~~~~ Segue 1 (19.0~$\pm$~0.6)\\ The number in parenthesis is $\log_{10}(L\times$\,GeV$^{-2}$cm$^{5}$), where $L$ is the astrophysical factor as defined in Eq. (\ref{jfactor}) and calculated in \cite{rouven, newsegue} \footnote{The astrophysical factor for Segue 1 listed here represents an updated value \cite{newsegue} that was not available at the time of writing of \cite{fsrpaper}; all plots and discussions in the following sections use this updated value.}; it should be noted that the uncertainties on these astrophysical factors are extremely large at present. The Sloan Digital Sky Survey has recently discovered many new dwarf galaxies, and since only a small region of the galactic neighborhood has been completely surveyed, several hundred more low-luminosity, dark matter dominated dwarf galaxies might still be discovered in the future. It would be straightforward to apply the analysis of this paper to any promising new dwarf that may be discovered, once the distribution of dark matter is mapped out to allow for at least an approximate determination of its $L$ factor. \textit{Atmospheric Cherenkov telescopes (ACTs) ---} The key parameters governing the sensitivity of an ACT in observations of dwarf galaxies are its effective area $A_{eff}$, energy resolution $\epsilon$, and energy threshold. There are several currently operational (eg. MAGIC, VERITAS) and near future (eg. CTA) ACTs relevant for indirect dark matter detection; we refer to \cite{fsrpaper} for more details, references, and individual key parameter values. Typically, current ACTs have $A_{eff}\sim10^9cm^2$ and $\epsilon\sim 0.15$, while future instruments are expected to reach $A_{eff}\sim10^{10}cm^2$ and $\epsilon\sim 0.10$. The typical instrumental energy threshold for ACT's is about 200 GeV. \textit{Previous observations and upper bounds ---} The dwarf galaxies mentioned above have been observed by ACTs without any positive detection, resulting in upper bounds on high energy gamma ray flux from dark matter annihilation or decay in these galaxies. Figure \ref{bounds}, which compares these experimental bounds with theoretical predictions from the three leptophilic models of interest in this paper, shows that the predictions are consistent with the observed null results within the uncertainties in the astrophysical factors; for more details on bounds from individual observations, the reader is referred to \cite{fsrpaper} and references therein. \begin{figure}[t] \centering \includegraphics[width=3.5in,height=1.8in]{bounds.eps} \caption{Comparison of experimental bounds with predictions from theory. The horizontal lines represent the bounds from experimental observations (\cite{fsrpaper} and references therein). The three vertical bars for each search are the corresponding predictions of models A (left bar), B (center), and C (right), using the astrophysical factors as listed in the text, with circles denoting central values.} \label{bounds} \end{figure} \textit{Backgrounds ---} Since the dwarf galaxies themselves are not expected to contain significant sources of hard gamma rays of astrophysical origin, the background can be effectively measured by looking at a region of the sky close to the dwarf (called the OFF region); subtracting the OFF region flux from the flux in the ON region (which contains the dwarf) eliminates the background up to statistical fluctuations. Astrophysical backgrounds can be estimated with standard extrapolations of charged lepton, hadron, and gamma ray spectra \cite{bergstrom, robust}; the hadronic background depends on the hadron rejection capabilities of the instrument ( see \cite{fsrpaper} for a more detailed treatment). Gamma rays from dark matter annihilation within the Milky Way also contribute to this background. In addition to FSR, these come from inverse Compton scattering (ICS) of starlight and CMB photons off energetic leptons from dark matter annihilation; we estimate the ICS contribution using the semi-analytic formalism in \cite{positrons}.

If the PAMELA, Fermi and HESS anomalies have their origin in leptophilic dark matter annihilation, current and near-future ACTs have an excellent chance of observing the accompanying final state radiation from dwarf galaxies. Unfortunately, lack of precise knowledge of the distribution of dark matter in the dwarfs makes the signal flux predictions highly uncertain. If a signal is observed, the measured gamma ray spectrum can likely be used to identify the correct annihilation channel and dark matter mass --- a general conclusion that holds for a range of signal strengths, dark matter masses, energy thresholds, and instrument parameters --- paving the way to a better understanding of the microscopic nature of dark matter. \vskip0.5cm \noindent{\large \bf Acknowledgments} \vskip0.3cm This research is supported by the U.S. National Science Foundation through grant PHY-0757868 and CAREER award PHY-0844667.

1012.4024

1012

1012.1747_arXiv.txt

{The advent of large instantaneous bandwidth receivers and high spectral resolution spectrometers on (sub-)millimeter telescopes has opened up the possibilities for unbiased spectral surveys. Because of the large amount of data they contain, any analysis of these surveys requires dedicated software tools. Here we present an extension of the widely used CLASS software that we developed to that purpose. This extension, named Weeds, allows for searches in atomic and molecular lines databases (e.g. JPL or CDMS) that may be accessed over the internet using a virtual observatory (VO) compliant protocol. The package permits a quick navigation across a spectral survey to search for lines of a given \edit{species}. Weeds is also capable of modeling a spectrum, as often needed for line identification. We expect that Weeds will be useful for analyzing and interpreting the spectral surveys that will be done with the HIFI instrument on board Herschel, but also observations carried-out with ground based millimeter and sub-millimeter telescopes and interferometers, such as \edit{IRAM-30m and Plateau de Bure, CARMA, SMA, eVLA, and ALMA.}}

A spectral survey consists in a series of spectra covering a significant spectral domain. At (sub-)millimeter wavelengths, a spectral survey typically covers several tens of GHz. Spectral surveys are generally \edit{referred} to as unbiased if they provide a complete coverage with a uniform sensitivity. As such, they allow for a complete census of the species emitting in that band, and sometimes for discovery of new interstellar species. In addition, because a given band often contains many transitions of the same species, the simultaneous analysis of all these lines provides stringent constraints on the physical conditions in the emitting gas, such as the density and temperature. Therefore spectral surveys are very useful for characterizing both the chemical composition and physical condition in the observed objects. Ever since the pioneering work of \citet{Johansson84}, who carried-out an unbiased spectral survey of the Orion~KL star-forming region and IRC~+10216 carbon-rich star between 72 and 91~GHz with the Onsala telescope, many spectral surveys have been carried-out at millimeter and sub-millimeter wavelengths using ground-based telescopes \citep[see][for a review]{Herbst09}. Because of the limited sensitivity of the instruments available at that time, early spectral surveys were targeted at bright star-forming regions, such as Orion~KL and Sgr~B2 in the millimeter range. Thanks to the increasing sensitivity of heterodyne receivers and the availability of sub-millimeter telescopes, these surveys were later extended to higher frequencies \cite[e.g.][]{Schilke97a,Schilke01,Comito05} and carried-out towards fainter young stellar objects \citep[e.g. NGC~1333~IRAS4 or IRAS16292-2422;][]{Blake94,vanDishoeck95,Blake95}. \edit{A few spectral surveys have been carried with millimeter and sub-millimeter interferometers, such as OVRO or the SMA \citep[e.g.][]{Blake96b,Beuther06}.} The HIFI instrument \citep{deGraauw10} onboard the Herschel space observatory \citep{Pilbratt10} now allows for a complete coverage of the almost unexplored 480-1250 and 1410-1910~GHz frequency bands. Its large spectral coverage -- up to 4~GHz instantaneous bandwidth -- and unprecedented sensitivity in this frequency range enable astronomers to carry-out spectral surveys over almost 1.5~THz down to the line confusion limit in a few tens of hours. The first spectral surveys with this instrument have already given spectacular results \citep{Bergin10,Ceccarelli10}. Among these, we can cite the richness of the Orion BN-KL spectrum observed at THz frequencies \citep[see Fig.~2][]{Bergin10} or the discovery of ND in IRAS16293-2422 \citep{Bacmann10}. Current \edit{developments} in \edit{(sub-)millimeter instruments} include an increase in the instantaneous bandwidth of the detection devices. During the past decade, the instantaneous bandwidth of tunable heterodyne receivers has increased by more than an order of magnitude, now routinely reaching $\sim$10 GHz. Other technologies (e.g. HEMT, FCRAO and IRAM) have already provided several tens of GHz, although it is still unclear whether the sensitivity of these receivers can match that of SIS receivers. This increase in bandwidth has been accomplished in parallel with the advent of digital spectrometers (autocorrelators, fast Fourier transform), the versatility of which allow the coverage of such bandwidth with a spectral resolution down to a few hundred kHz. As a result, unbiased spectral surveys of the 3 mm atmospheric window ($\nu = 80 - 117$~GHz) can be done with the IRAM-30m telescope in $\sim$10~hours, with a 2~mK noise at 1$\sigma$ in 2~MHz ($\sim$6~km/s) spectral channels. The ALMA interferometer will also permits coverage of large frequency windows, providing spectral cubes with up to 8~GHz bandwidth \citep{Wooten08}. \edit{Thanks to its sensitivity, this instrument will allow, in its compact configuration, line surveys to be carried-out down to the confusion limit toward a large number of sources}. Spectral surveys are thus still in their infancy and will very likely become routine observing modes in the coming years. Spectral surveys covering large frequency bands require specific tools to be analyzed efficiently. In this article, we present a software that is intended for the analysis of spectral surveys. In \S~\ref{sec:spectr-surv-analys}, we briefly describe how such surveys are analyzed. In \S~\ref{sec:weeds-design-impl} we detail how our software was designed and implemented to carried-out such an analysis. Finally \S~\ref{sec:concl-prosp} concludes this article and discuss future developments.

\label{sec:concl-prosp} We have presented an extension of the CLASS data reduction software for analyzing spectral surveys. This extension allows the user to make queries in spectral line databases using a VO compliant protocol. It also allows the user to quickly search for the various transitions of a given specie. Finally it can compute model predictions at the LTE, as often needed to identify lines in spectra close to the confusion limit. Weeds has already been successfully used to analyze part of the IRAS~16293-2422 survey obtained with \emph{Herschel-HIFI} \citep{Bacmann10,HilyBlant10}, and we expect that it will be useful for future spectral surveys with this instrument as well. We think that it will become a standard tool for analyzing spectral surveys obtained with single dish ground based telescopes such as the IRAM-30m. Yet, Weeds is not limited to the analysis of single dish observations. \edit{It may be used to analyze spectral surveys obtained with interferometers as well, such as the IRAM Plateau de Bure, CARMA, the SMA, and the upcoming ALMA and eVLA interferometers. In fact, since} Weeds is written in Python, it could be used from the Python based CASA software, that will be used by the eVLA and ALMA. However, analyzing ALMA data will be challenging, because these data will consist in large spectral cubes, i.e. essentially a spectral survey on large number of pixels. In fact, doing such an analysis by hand, i.e. identifying the various lines/species on each spectrum of map is probably impossible; this will require some automatic fitting tools to extract the relevant information (column densities and excitation temperature of the various species) as a function of position. Such tools require efficient minimization algorithms to fit a model with a large number of free parameters to the data. The development of such tools is already in progress \edit{(e.g. in XCLASS using the MAGIX minimization framework), and implementing these automatic fitting tools in Weeds would be desirable in the future.}

1012.1747

1012

1012.4354_arXiv.txt

In turbulent dynamos the production of large-scale magnetic fields is accompanied by a separation of magnetic helicity in scale. The large- and small-scale parts increase in magnitude. The small-scale part can eventually work against the dynamo and quench it, especially at high magnetic Reynolds numbers. A one-dimensional mean-field model of a dynamo is presented where diffusive magnetic helicity fluxes within the domain are important. It turns out that this effect helps to alleviate the quenching. Here we show that internal magnetic helicity fluxes, even within one hemisphere, can be important for alleviating catastrophic quenching.

1012.4354

1012

1012.3268_arXiv.txt

A mapping study of IRAS 05553+1631 was performed with $^{12}$CO J=3-2 and $^{13}$CO J=2-1 lines observed by the KOSMA 3 m-telescope. A core with a size of 0.65 pc and with a LTE mass of 120 M$_\odot$ was defined by the mapping with $^{13}$CO J=2-1 line. We have identified a bipolar outflow with $^{12}$CO J=3-2. For accuracy in the calculation of outflow parameters, overcoming the projection effect is important. We propose a new method to directly calculate the inclination-angle $\theta$. We establish two basic equations with the help of outflow contour diagram and finally obtain the "angle function" and the "angle equation" to derive $\theta$. We apply our method to the outflow of IRAS 05553+1631, finding that $\theta_{blue}$ is 73$^\circ$ and $\theta_{red}$ is 78$^\circ$. Compared to the parameters initially estimated under an assumption of 45$^\circ$ inclination-angle, the newly derived parameters are changed with different factors. For instance, the timescales for the blue and the red lobes are reduced by 0.31 and 0.21, respectively. Larger influences apply to mechanical luminosity, driving force, and mass-loss rate. The comparisons between parameters before and after the correction show that the effect of the inclination-angle cannot be neglected.

Massive star formation (MSF) has attracted much attention. It has enormous impact on the natal interstellar medium (ISM) and on the evolution of Galaxy. MSF is believed to originate in dense molecular cores (DMCs) that can be traced by CO and its isotopologues. With the help of mapping, we can reveal the structure of molecular cores and investigate the MSF activity taking place within them. High velocity outflows have also been intensively studied. First uncovered in 1976 \citep{b19,b8}, molecular outflows toward massive young stellar objects (YSOs) have attracted much attention \citep{b16}. Perhaps tracing the earliest stage \citep{b9} of star formation, molecular outflows are critical in the debate about two mechanisms for MSF: massive stars forming through accretion-disk-outflow or via collision-coalescence \citep{b22,b23}. The properties of outflows reveal the mass-loss phase before the main sequence. What we observe in the sky is not the real outflow but its two-dimensional projection. Therefore, it is essential to know the inclination-angle $\theta$ between the outflow axis and our line-of-sight. For instance, it will introduce a factor of 1/cos$\theta$ to the momentum. So far the determination of the inclination-angle has been somewhat neglected though some previous researchers have made a few attempts by modeling \citep{b3,b10}. Until now, authors usually took assumptions for the inclination-angle. For instance, \citet{b7} hypothesized 45$^\circ$ while \citet{b6} supposed 60$^\circ$. The outflow parameters derived from these assumptions can only be meaningful statistically, but they are not accurate. In addition, work has been done to study the collimation of outflows \citep{b1,b15}. However, the projection effect should also be included. In order to directly calculate the inclination-angle and to eliminate or reduce outflow parameter uncertainties, we propose a new method in this paper. In the next section, we will describe our observations of IRAS 05553+1631. In section \ref{sec:results}, the results will be presented. We will introduce our method and discuss some of its properties in section \ref{sec:discussion}. Section \ref{sec:summary} is a brief summary of our work.

\label{sec:discussion} \subsection{Core} Table \ref{tab:core} shows that the core has a size $\sim$0.65 pc and that M$_{LTE}$ is equal to 120 $M_{\odot}$. The $^{13}$CO line width is 2.5 km s$^{-1}$ which is larger than those of low-mass cores \citep{b11}. These results suggest that the core is massive. According to virial theory, a core is gravitationally bound when its mass is larger than its virial mass. In our results, this is not satisfied, which is possibly due to the [H$_2$/$^{13}$CO] ratio adopted or the optically thick assumption for $^{13}$CO J=2-1 emission; alternatively, the core may be unbound. \subsection{Outflow}\label{sec:outflow} In our calculation for outflow parameters, 45$^\circ$ was assumed as the inclination-angle. Actually, almost all the outflow parameters depend on the angle. For instance, the velocity of outflow gas is the projection along our line-of-sight. Therefore a factor of 1/cos$\theta$ should be introduced. The same goes for momentum, timescale, etc., but sometimes with different factors. The inclination-angle also affects the collimation factor. Thus, accurate determination of $\theta$ is important. Unfortunately, until now, people assume a value for $\theta$ \citep{b7,b6,b18}. In the following, we propose a new method to improve this situation. \subsubsection{The inclination-angle $\theta$} Our method to derive $\theta$ requires knowledge of the outflow geometry, so before the calculation, it is necessary to discuss something about the shapes of outflows. Generally, without impact from the surrounding materials, the molecular outflows are likely to be axially symmetric. This could be inferred from the symmetry of the outflow contour diagrams (\citet{b7}, Figure 2; \citet{b17}, Figure 5c; \citet{b1}, Figure 13 NGC 2071) and from modeling \citep{b20,b10,b21}. In our results, the P-V diagram (Figure \ref{fig:12CO}) and the outflow contours (Figure \ref{fig:outflow}) show good symmetry, especially the blue lobe. In addition, the ellipse-like outflow contours indicate a projection of a conical outflow (see Section \ref{sec:more}). Since a cone is an ideal form for an axially symmetric outflow (see Section \ref{sec:comparison}, Paragraph 2), we assume that the shape of the outflow is a cone. We also assume that high velocity components form cone-shell like structure whose emission intensity decreases from inside to outside. The cone could be called a "cone onion". When we observe, what we see is the outflow projection on the sky. Thus, if the outflow axis is at an angle $\theta$ to our line-of-sight, the contours will appear as ellipses (Figure \ref{fig:cone}; the contours could also be parabola or hyperbola, see section \ref{sec:more}). Figure \ref{fig:model} presents the calculation of the inclination-angle $\theta$. In Figure \ref{fig:model}(a), one can see the driving centre and an elliptical outflow contour. From Figure \ref{fig:model}(b), a relation of the four angles: $\beta_1$, $\beta_2$, $\theta_1$, and $\theta$ can be expressed by equation (\ref{equ:4angle}) \begin{equation}\label{equ:4angle} \frac{\tan\theta_1}{\tan\theta}=\frac{\tan\beta_1}{\tan\beta_2}\approx\frac{\beta_1}{\beta_2} \end{equation} where $\beta_1$ and $\beta_2$ are the angle distances in the contour-diagram ($\beta_1$ accounts for the angle between driving centre and point D; $\beta_2$ is the angle between driving centre and axis point A). $\beta_1$ and $\beta_2$ are usually very small (about several arcmin), thus the small angle approximation is valid. $\theta_1$ and $\theta$ are unknown. Then, we have to find another equation for either $\theta_1$ or $\theta$. Notice that the length of line AB does not change as a function of $\theta$. In Figure \ref{fig:model}(a), we have \begin{equation}\label{equ:a} \overline{AB}=\overline{AO}\tan\alpha \end{equation} where $\alpha$ is the angle between OA and OB. Overlines indicate length. If the outflow contour is not symmetric about line AO, we take $\alpha$ as half the angle $\angle BOC$. We have already assumed the shape of cone, so the opening angle $\theta$-$\theta_1$ is the same for the entire cone. Therefore in Figure \ref{fig:model}(b) we have \begin{equation}\label{equ:b} \overline{AO_1}=\overline{AO_2}\sin\theta \end{equation} \begin{equation}\label{equ:c} \overline{AB}=\overline{AO_2}\tan(\theta-\theta_1) \end{equation} where O$_1$ and O$_2$ are the same points as O in the plane of the sky. $\overline{AO_1}$ is equal to $\overline{AO}$. Thus, combining equations (\ref{equ:a}), (\ref{equ:b}), (\ref{equ:c}), and substitute $\overline{AO_1}$ with $\overline{AO}$ we have \begin{equation}\label{equ:alpha} \tan\alpha=\frac{\tan(\theta-\theta_1)}{\sin\theta} \end{equation} Putting together equation (\ref{equ:4angle}) and equation (\ref{equ:alpha}), we can solve the inclination-angle $\theta$ by eliminating $\theta_1$. Since $\beta_1$, $\beta_2$, and $\alpha$ can be derived from the outflow contour, we try to further manipulate the two equations to obtain a better expression. We define the "angle constants" \begin{equation} P=\frac{\beta_1}{\beta_2} \end{equation} and \begin{equation} Q=\tan\alpha \end{equation} Combining equations (\ref{equ:4angle}) and (\ref{equ:alpha}) and eliminating $\theta_1$, we have \begin{equation} \frac{1-P}{Q}\frac{1}{\cos\theta}-P\tan\theta\tan\theta=1 \end{equation} Defining the "angle function" \begin{equation}\label{angle function} A(\theta)=\frac{1-P}{Q}\frac{1}{\cos\theta}-P\tan\theta\tan\theta, \end{equation} after determining P and Q, we obtain the inclination-angle by solving the "angle equation" \begin{equation}\label{angle equation} A(\theta)=1 \end{equation} \subsubsection{Comparison}\label{sec:comparison} We utilize our method on IRAS 05553+1631. Figure \ref{fig:theta} shows the derivation process. In general, the shapes and positions in Figure \ref{fig:theta} are similar to those in our model (Figure \ref{fig:model}(a)). Table \ref{tab:comparison} presents the comparison between newly derived and old outflow parameters. We find the inclination-angles to be 73$^\circ$ and 78$^\circ$ for blue and red lobes, respectively. This suggests that the two outflow lobes are in good alignment. In Table \ref{tab:comparison}, compared with the former results, the newly derived parameters have rather large corrections. For the blue lobe, the momentum is enlarged by a factor of (cos45$^\circ$/cos73$^\circ\approx$) 2.4, as the velocity is corrected by 1/cos$\theta$. The correction for energy is proportional to that of P$^2$ which is about 5.8. The timescale is reduced by a factor of (cot73$^\circ$/cot45$^\circ\approx$) 0.31. The correction factors for mechanical luminosity, driving force, and mass-loss rate are 19, 7.7, and 7.7, respectively. For the red lobe, the analysis is similar, but the factors are different. These results suggest that the correction from inclination-angle cannot be neglected and our method is practical. From the process of our method one can see that the cone assumption ensure the viability of equation (\ref{equ:alpha}). If not, the equation would be an approximation. Let $\theta_0$=$\theta$-$\theta_1$, the Taylor expansion of tan$\theta$ around $\theta_0$ is (keeping the first order) tan$\theta$=tan$\theta_0$+$\frac{\Delta\theta}{(cos\theta_0)^2}$. For the blue outflow lobe in our case, tan$\theta_0$=0.143, if $\Delta\theta$=1$^\circ$ (0.017 radians), then the second term of the Taylor expansion would be $\sim$0.018, the uncertainty is about 12 per cent. For the red outflow lobe, the uncertainty is about 4 per cent. \subsubsection{More properties}\label{sec:more} Geometrically, when a cone is cut by a plane, the resulting curve can be an ellipse (including a circle), a parabola, or a hyperbola, depending on the angle between the cone-axis and the plane. For an ellipse, the larger the inclination-angle is, the larger its eccentricity will be, and vice versa. When $\theta$ approaches 90$^\circ$ we obtain a parabola or a hyperbola. Theoretically, all of the curves can be utilized for the calculation. But the ellipse has the advantage that it is usually much more conspicuous and easy to manipulate. The parabola and the hyperbola, on the contrary, can be confusing because of the large inclination-angle. The outflow edge from the driving source has very weak emission and the contours exhibit a fan-like structure. The centre and the curve will be ambiguous. The opposite extreme is when the inclination-angle becomes zero and the ellipse changes to a circle. Fortunately, the angle equation (\ref{angle equation}) still works. In the angle function (\ref{angle function}), the coefficient of the first term $\frac{1-P}{Q}$ is ($\frac{1-\frac{\overline{OD}}{\overline{OA}}}{\frac{\overline{AB}}{\overline{OA}}}$=$\frac{\overline{OA}-\overline{OD}}{\overline{AB}}$=$\frac{\overline{AD}}{\overline{AB}}$). As $\theta$ approaches zero, the projection shape becomes a circle and $\frac{\overline{AD}}{\overline{AB}}\rightarrow$1. Meanwhile, cos$\theta$ $\rightarrow$1 and the second term vanishes. Thus A($\theta\rightarrow$0)$\rightarrow$1. The angle equation (\ref{angle equation}) is tenable in this extreme case. However, when $\theta$ approaches zero, use of this method could be problematic if the resolution is low. Another property is the collimation. In our view, it can be manifested by the opening angle $\theta$-$\theta_1$. Larger $\theta$-$\theta_1$ means lower collimation, and vice versa. Thus it is reasonable to define a collimation factor as cot($\theta$-$\theta_1$) (=(tan$\alpha$sin$\theta$)$^{-1}$). In IRAS 05553+1631, the factors are 7.0 and 1.6 for blue and red lobes, respectively. One thing should be mentioned, the ellipses we used for IRAS 05553+1631 are 90 per cent contours (Figure \ref{fig:theta}) as stronger emission could reduce the errors. If we chose the 50 per cent contours, the factor would change. Additionally, the method highly depends on the accuracy with which the driving centre is located. One question is the identification, and the second is the spatial resolution. The identification involves personal judgement, which is accompanied by considerable uncertainty in some cases. \citet{b4} showed that 6.7 GHz methanol maser and relevant 24 $\mu m$ emission usually coincide with each other and they are good tracers to the driving centre, this may offer much help for the identification. The resolution is usually low for single-dish observations. Using high-resolution telescopes and interferometers can improve this. In the angle function (\ref{angle function}), there are two parameters P and Q. Since we can always choose an ellipse excluding the driving centre, it is reasonable to assume 0$<$P$<$1. Figure \ref{fig:Q} illustrates the value of Q, which shows the projection on x-z plane. Lines AO$_1$ and AD' represent the projections of two conditions for the plane of the sky: (1) AO$_1$ perpendicular to O$_2$E (solid, hereafter named c1); (2) AD' not perpendicular to O$_2$E (dashed, hereafter named c2). Condition c2 is denoted by prime. $\overline{AE}$ equals the $\overline{AB}$ of Figure \ref{fig:model}(a). Thus, it is easy to see that the maximum of Q is $\frac{\overline{AE}}{\overline{AO_1}}$ which is just 1/cos($\theta$-$\theta_1$). Since ($\theta$-$\theta_1$) cannot be large and an angle of 60$^\circ$ can only bring Q=2, we consider 0$<$Q$<$2. When P=0, the second term of the angle function (\ref{angle function}) vanishes and $\theta$=arccos(1/Q). In order to have solution, Q must be larger than 1. In fact, in this case Q reaches its maximum (see Figure \ref{fig:Q}). The result is 0$<\theta<$71$^\circ$. When P=1, there is no solution. Figure \ref{fig:curves} shows the variation of the angle function as a function of P and Q. With a fixed Q, the function decreases as P increases. With a fixed P, the function decreases when Q increases. Usually within (0$^\circ$,90$^\circ$), there is one solution. When Q approaches 1 (Figure \ref{fig:curves}(d)), in order to have solution, P must be closer to 0. When P approaches 1, the same happens to Q. The variation of angle function is sensitive to the values of P and Q. We mapped IRAS 05553+1631 with $^{12}$CO J=3-2 and $^{13}$CO J=2-1 lines. A core was identified from $^{13}$CO J=2-1 observations. It has a size of 0.65 pc and LTE mass of 120 M$_{\odot}$ which is lower than the virial mass of 850 M$_{\odot}$. $^{12}$CO J=3-2 mapping revealed a bipolar outflow. Its parameters were initially estimated under the assumption of a 45$^\circ$ inclination-angle. A new method to directly calculate the inclination-angle $\theta$ was proposed, and was utilized for the bipolar outflow of IRAS 05553+1631. We found that $\theta_{blue}$ is 73$^\circ$ and $\theta_{red}$ is 78$^\circ$. Parameters with the new $\theta$s were compared with the former ones. For the blue lobe, the momentum was enlarged from 82 M$_\odot$~km~s$^{-1}$ to 200 M$_\odot$~km~s$^{-1}$ by a factor of 2.4 while the timescale was reduced from 8.8$\times$10$^4$ yrs to 2.7$\times$10$^4$ yrs by a factor of 0.31. The enlarging factors for energy, mechanical luminosity, driving force, and mass-loss rate are 5.8, 19, 7.7, and 7.7, respectively. For the red lobe, the momentum was enlarged from 11 M$_\odot$~km~s$^{-1}$ to 36 M$_\odot$~km~s$^{-1}$ by a factor of 3.4 while the timescale was reduced from 6.3$\times$10$^4$ yrs to 1.3$\times$10$^4$ yrs by a factor of 0.21. The enlarging factors for energy, mechanical luminosity, driving force, and mass-loss rate are 12, 55, 16, and 16, respectively. The results show that a selection of parameters were influenced by the inclination-angle $\theta$.

1012.3268

1012

1012.1092_arXiv.txt

Over cosmic time, galaxies grow through the hierarchical merging of smaller galaxies. However, the bright region of the galaxy luminosity function is incompatible with the simplest version of hierarchical merging, and it is believed that feedback from the central black hole in the host galaxies reduces the number of bright galaxies and regulates the co-evolution of black hole and host galaxy. Numerous simulations of galaxy evolution have attempted to include the physical effects of such feedback with a resolution usually exceeding a kiloparsec. However, interactions between jets and the interstellar medium involve processes occurring on less than kiloparsec scales. In order to further the understanding of processes occurring on such scales, we present a suite of simulations of relativistic jets interacting with a fractal two-phase interstellar medium with a resolution of two parsecs and a largest scale of one kiloparsec. The transfer of energy and momentum to the interstellar medium is considerable, and we find that jets with powers in the range of $10^{43}$--$10^{46} \ergs$ can inhibit star formation through the dispersal of dense gas in the galaxy core. We determine the effectiveness of this process as a function of the ratio of the jet power to the Eddington luminosity of the black hole, the pressure of the interstellar medium and the porosity of the dense gas.

\label{s:intro} It is widely believed that feedback from active galactic nuclei (AGN) during the epoch of galaxy formation is required to explain the relation between black hole and bulge mass/velocity dispersion \citep{magorrian98a,gebhardt00a,tremaine02a}, the deficit of bright galaxies in the galaxy luminosity function \citep{cole01a,norberg02a,huang03a}, and the completion of star formation in massive galaxies at epochs of redshift $z \lesssim 2$ \citep{shaver96a,madau96a,bender99a}. It is envisaged that either radiation or outflows from a galactic nucleus impedes the infall of star-forming gas once the central black hole grows to a critical size. Accordingly, \citet{silk98a}, \citet{fabian99a}, \citet{king05a}, and others, have appealed to the physics of either energy-driven or momentum-driven bubbles in order to explain the relationships between the mass of the black hole and the parameters of the host galaxy. In order to model the galaxy luminosity function, \citet{croton06a} have utilized semi-analytic models based on the output of the Millennium Simulation, incorporating \lq\lq{}radio-mode\rq\rq{} feedback, coupled with a prescription for the accretion rate into the center of each evolving galaxy. Their feedback prescription is motivated by the well-documented evidence for the effect of radio galaxies on \lq\lq{}cooling flow\rq\rq{} galaxies \citep{fabian03a,mcnamara05a}. There is also a growing literature on cosmological simulations in a \LCDM{} cosmogony involving both dark matter and gas dynamics, which incorporate feedback from both supernovae (SN) and black holes, and which test scenarios of galaxy merging and growth \citep{springel03a,springel03b,booth09a,booth09b,schaye10a}. In simulations using the Gadget-2 code \citep{springel05a} the total number of SPH particles representing the baryonic component exceeds 250 million, and the effective spatial dynamic range is an equally impressive $10^5$ per dimension. Nevertheless, the best spatial resolution is about 2 kpc and does not resolve the spatial scales where important dynamical processes occur. This is highlighted by the prescriptions for black hole growth and feedback described, for example in the work by \citet{booth09b}. Accretion is described in terms of the Bondi-Hoyle accretion rate multiplied by a factor, which can be as large as $10^2$. The rationale for this approach is that at sub-grid scales the density would be larger, and the real accretion rate would be appreciably higher. However, the higher densities and the consequent cooling and fragmentation on sub-grid kpc scales would create a multi-phase interstellar medium (ISM), the physics of which is not satisfactorily captured by the simulations. In particular, such a medium is porous, and the interaction between jets and the dense, potentially star-forming clouds of gas is complex, with radio-emitting plasma being able to channel through holes in the density distribution, rather than isotropically impacting a smooth distribution of dense gas as shown in previous work \citep{sutherland07a}. Given that black-hole driven feedback occurs in bright galaxies, there does not appear to be a consensus on the type of feedback: How much AGN power is involved, and does the feedback involve radiative or mechanical processes or both? In their radio-mode model, \citet{croton06a} attribute feedback to radio galaxies accreting at rates well below Eddington, and for typical ellipticals, this means low-powered Fanaroff-Riley Class I radio galaxies are the primary drivers of feedback. In \citet{fabian99a} the momentum for dispersing the circumnuclear gas comes from a quasar wind; in \citet{king05a} the momentum of the outflow is provided by an Eddington-limited radiation-driven wind; in SPH simulations the accretion rate, which is the ultimate power source for an outflow, can approach Eddington values \citep{booth09a}. There is also an issue of what class of AGN actually drives black hole feedback. The models by \citet{croton06a} invoke low-powered radio galaxies. Observers, however, have focused on powerful radio galaxies, mainly at $z\gtrsim 2$; in these galaxies there is evidence for substantial outflows of line-emitting gas and neutral gas driven by the radio jets \citep{nesvadba09b, morganti10a}. It is possible that there is a role for radio galaxies with a range of powers: Powerful radio galaxies may be responsible for the establishment of the stellar mass at $z\gtrsim 2$, and less powerful sources may be responsible for the maintenance of the stellar content through the inhibition of cooling flows \citep{nulsen09a}. We also note that radio galaxies are mainly relevant to the elliptical galaxy population, and that the separate luminosity functions for early and late-type galaxies \citep{huang03a} indicate the requirement for feedback in both populations. In this paper we consider the potential role of radio galaxies in AGN feedback and address the following questions: (1) What jet power is required for the radio galaxy phase to have an important effect on inhibiting star formation in a given host, and (2) Is the range of radio powers broad enough that radio galaxies could affect the entire distribution of bright ellipticals? In this paper we present progress in answering these questions through simulations with a resolution of $2\parsec$ per computational cell. These simulations confront the sub-grid physics that current large scale SPH simulations do not address. In particular, we consider the effect of powerful relativistic jets on a two-phase ISM consisting of hot gas, in which is embedded a dense porous phase of warm gas. These simulations extend the simulations described in \citep{sutherland07a}, in which a jet with a kinetic power of $3 \times 10^{43} \ergs$ propagates through an inhomogeneous medium in the form of an almost Keplerian fractal disk. It is evident from that simulation that in the geometry considered, jets of that power could not exert enough impact on the clouds to disperse them, and that the jets would not have a important effect on star formation in the core of the host galaxy. In the present simulations we consider jets with powers ranging from $10^{43}$ to $10^{46} \ergs$ propagating through a two-phase medium, in which the dense clouds are spherically distributed throughout a region $1\kpc$ in diameter. These initial data are meant to describe a typical protogalaxy, in which dense gas has accumulated in the core. The fractal distribution of the dense gas enables us to directly examine the effect of porosity on the evolution of potentially star-forming clouds. In the following sections we describe the parameters of the simulations in more detail and then discuss our results.

\label{s:discussion} The suite of simulations presented in this paper strongly reinforce the importance of inhomogeneity in the consideration of jet interactions with the ISM \citep{saxton05a,sutherland07a}. Inhomogeneity has several important consequences: (1) It affects the early morphology of the radio source as a result of the interaction of the jet and lobe with the obstructing clouds. (2) The radio source affects a much larger volume of the host galaxy because of the channeling of the jet flow in different directions. (3) The shocked jet clouds are left in the wake of the non-thermal plasma and, as first noted by \citet{sutherland07a}, would continue to emit shock excited line emission with the shocks driven by the high pressure gas in the non-thermal cocoon; this shock-excited emission is in addition to the emission that may be driven by photoionization by the nucleus. (4) The porosity of dense gas determines the ease with which a jet in a given host can disperse this gas, which is a potential source of new stars. In powerful sources, higher porosity gas is less easily dispersed but this trend is not evident for lower-powered ($\sim 10^{44} \ergs$) jets. (5) Jets of all powers can exert a considerable feedback effect on their host galaxies, although lower-powered jets only play a role in the lower velocity dispersion hosts. Brighter galaxies require more powerful jets to disperse dense clouds. (6) The efficiency of transfer of kinetic energy from the jet to the dense gas is high. (7) The efficiency of transfer of momentum to the clouds is also high with mechanical advantages considerably exceeding unity. Inhomogeneity is therefore crucial when considering AGN feedback on the kiloparsec scale, both for the interpretation of radio and optical emission-line morphology in radio sources, which may be generating feedback, and for incorporating the effect of jet-mediated feedback on host galaxies of different size into large scale simulations, in which the resolution $\gtrsim 1 \kpc$. An important conclusion from these simulations is that jets with Eddington efficiency $\eta \lesssim 10^{-4}$ are unlikely to have an effect on evolving galaxies when the pre-starforming gas exists in the form of clouds, which are relatively dense and cool compared to the hot ISM. This critical value of $\eta$ is relevant for clouds with a high filling factor of 0.42 and a value of $p/k = 10^6$ for the ISM. We have shown that when the filling factor decreases or the pressure of the ISM increases the critical value of $\eta$ increases, and values of $\etacrit \sim 10^{-3} - 10^{-2}$ are not unrealistic. The precise values of $\etacrit$ will have to await further higher resolution simulations with high porosity dense gas.

1012.1092

1012

1012.4481_arXiv.txt

The relation of the solar surface magnetic field with mesogranular cells is studied using high spatial ($\approx 100$ km) and temporal ($\approx 30$ sec) resolution data obtained with the IMaX instrument aboard SUNRISE. First, mesogranular cells are identified using Lagrange tracers ({\it corks}) based on horizontal velocity fields obtained through Local Correlation Tracking. After $\approx 20$ min of integration, the tracers delineate a sharp mesogranular network with lanes of width below about $280$ km. The preferential location of magnetic elements in mesogranular cells is tested quantitatively. Roughly $85\%$ of pixels with magnetic field higher than $100$ G are located in the near neighborhood of mesogranular lanes. Magnetic flux is therefore concentrated in mesogranular lanes rather than intergranular ones. Secondly, magnetic field extrapolations are performed to obtain field lines anchored in the observed flux elements. This analysis, therefore, is independent of the horizontal flows determined in the first part. A probability density function (PDF) is calculated for the distribution of distances between the footpoints of individual magnetic field lines. The PDF has an exponential shape at scales between $1$ and $10$ Mm, with a constant characteristic decay distance, indicating the absence of preferred convection scales in the mesogranular range. Our results support the view that mesogranulation is not an intrinsic convective scale (in the sense that it is not a primary energy-injection scale of solar convection), but also give quantitative confirmation that, nevertheless, the magnetic elements are preferentially found along mesogranular lanes.

\label{sec:introduction} Mesogranulation was historically introduced as a prominent scale imprinted on the horizontal photospheric flows calculated through local correlation tracking (LCT) of intensity images \citep{november1981, simon1988, brandt1988, brandt1991, muller1992, roudier1998, roudier1999, shine2000, leitzinger2005}: both the pattern of positive and negative divergence of that flow as well as the time evolution of Lagrange tracers moving in it revealed cells with sizes between, say, $5$ and $10$ arcsec. Much debate ensued concerning whether the mesogranular flow patterns correspond to actual convection cells in that range of sizes rather than, e.g., to simple granule associations which persist in time \citep[see, e.g.][]{cattaneo2001, roudier2003, roudier2004, nordlund2009, matloch2009, matloch2010}. Independent hints for the existence of a convective flow operating on those scales are therefore of importance. The study of the surface distribution of magnetic elements can provide such hints; as a minimum, it can constitute an alternative avenue, independent of the inaccuracies of LCT methods, to determine the properties of the mesogranular patterns. Such an approach has been employed by \citet{dominguezetal2003, dominguez2003, sanchez_almeida2003} using ground-based data and by \citet{roudier2009} and \citet{ishikawa2010} using satellite observations. In those papers, visual evidence was obtained that there is an association between magnetic flux structures (flux elements, transient horizontal fields) and the mesogranular pattern obtained through the study of horizontal flows. Yet, detailed quantitative studies and statistics that could put such an association on a firmer basis are still missing. The aim of this letter is to obtain quantitative information of the relation between photospheric magnetic flux distributions and mesogranular scales by using the observations of unprecedented quality provided by the Imaging Magnetograph eXperiment (IMaX; \citealt{martinezpillet_etal_2010}) aboard the SUNRISE balloon-borne observatory \citep{barthol_etal_2010, solanki_etal_2010}. IMaX provides time series of virtually seeing-free high spatial resolution ($0.15$ arcsec) images and magnetograms that constitute an ideal data set for the study of photospheric magnetism. Using the Sunrise/IMaX data, we combine information from (a) the velocity field gained through Local Correlation Tracking of intensity images; (b) the spatial patterns provided by the magnetograms; and (c) the field line structure obtained through extrapolation of the magnetogram data, to gain quantitative information concerning patterns at intermediate scales between granulation and supergranulation. \section {Data}\label{sec:data} \begin{figure} % \includegraphics[width=0.5\textwidth]{fig1.pdf} \caption{Top: normalized continuum intensity near the spectral line Fe I 5250.2 \AA. Bottom: Vertical component of the magnetic field vector ($B_z$) retrieved from inversions and clipped at $\pm 50$ G. The maps are taken from the first observed time series.} \label{fig1} \end{figure} For this study we use sequences of images recorded with IMaX near the solar disk center on 2009 June 9. Images were taken at five wavelengths along the profile of the magnetic-sensitive FeI $5250.2$ \AA\ line located at $\pm 80$ m\AA, $\pm 40$ m\AA\ from line center, and continuum at $+227$ m\AA. The estimated circular polarization noise is $5\times 10^{-4}$ in units of the continuum wavelength for non-reconstructed data and $3$ times larger for the reconstructed one. IMaX has a spectral resolution of $85$ m\AA\ and a spatial resolution of $\sim 100$ km \citep{martinezpillet_etal_2010}. The reduction procedure produces time series of images with a cadence of $33.25$ s, spatial sampling of $39.9$ km and a field-of-view (FOV) of $32 \times 32$ Mm$^2$. We use two time series, the first one comprising $42$ snapshots ($23$ minutes) and the second one $58$ snapshots ($32$ minutes). Magnetograms are derived from inversions of the observed Stokes parameters using the SIR code \citep{ruiz_cobo1992}. We call these magnetograms the {\it reconstructed data}. In this letter we mainly use the reconstructed data except in the last section (Sec.~\ref{sec:extrapolation}) where, in addition, we use so-called {\it calibrated data}, i.e., magnetograms obtained using a proportionality law between Stokes-V (from non-reconstructed data) and $B_z$ \citep[for details of this method, see][]{martinezpillet_etal_2010}. The intensity maps are taken from the reconstructed data. Figure~\ref{fig1} displays a map of normalized continuum intensity (top) and the corresponding longitudinal magnetogram (bottom), the latter showing many internetwork flux concentrations alongside stronger flux elements probably belonging to the network.

An important open question in solar physics is the precise nature of the convection scales with size and duration above the granular values. While both granulation and supergranulation yield velocity field patterns that have been observed at the surface, mesogranules have been detected only through indirect proxies, like tracking of intensity patterns, which do not provide reliable evidence of underlying convection cells in that range of sizes and durations. There is an ongoing debate on whether there is a continuum of sizes for the convection cells on scales above granular, possibly with self-similar properties and with no particular scale being singled out within that range \citep{nordlund2009}, or whether the mesogranular scales are just the result of a collective interaction between families of granules. The magnetic field can provide an alternative and more direct avenue to explore convective patterns since the magnetic flux can be measured directly using Stokes polarimetry techniques. The magnetic elements appear with a broad spectrum of flux densities \citep{orozco_suarez_etal_hinode_magnetic_elements_2007, khomenko_etal_2003} and can also be detected through G-Band or Ca-II bright points \citep[see][]{de_wijn_etal_2008, sanchezalmeida_etal_2010} and therefore can provide important clues concerning the nature of the flows underlying the mesogranular scales as well as about the magnetic elements themselves. The high spatial and temporal cadences of the IMaX data allow us to try a few different, complementary studies of the relation between magnetic field and mesogranular flows. First, we have carried out Lagrange tracing of mass elements following horizontal flow fields obtained through LCT of intensity maps; we have obtained a number of improvements compared with traditional cork maps (faster development of the mesogranular lanes, no need for time averages). We have also obtained an upper bound for the width of the mesogranular lanes of some $280$ km. Second, we have provided quantitative measures for the association between magnetic elements and mesogranular lanes. The large majority ($85$ \%) of the magnetic elements with flux density above $\sim 100$ G are found in $120$-km neighborhoods of mesogranular lanes with $\rhocork > 8$ and about $80$ \% of flux elements above $30$ G are located near mesogranular lanes with $\rhocork > 2$. Our results indicate a good coupling between the flow field and the magnetic elements, suggesting that the evolution of the latter is mostly kinematic. Third, we have considered the connectivity between magnetic elements and, in particular, the distance between footpoints of field lines anchored in the photosphere. The probability density function of such distances shows that there is abundant connectivity on mesogranular scales; it also shows that the distribution is basically featureless in those scales, with only one characteristic value, the slope of the distribution, equal to $(1.7\, \hbox{Mm})^{-1}$. Our results concerning statistics of separations of field line footpoints suggest that there is no intrinsic scale of convection in the mesogranular range. This may only mean that there is no mechanism for direct injection of energy into convection on those scales, but does not rule out the existence of convection cells with those sizes, e.g., through nonlinear interactions of cells at other scales (as in a turbulent cascade) or through the interaction of thermal downflows \citep{rast2003}. \vskip 5mm

1012.4481

1012

1012.0577_arXiv.txt

M dwarfs are known to flare on timescales from minutes to hours, with flux increases of several magnitudes in the blue/near-UV. These frequent, powerful events, which are caused by magnetic reconnection, will have a strong observational signature in large, time-domain surveys. The radiation and particle fluxes from flares may also exert a significant influence on the atmospheres of orbiting planets, and affect their habitability. We present a statistical model of flaring M dwarfs in the Galaxy that allows us to predict the observed flare rate along a given line of sight for a particular survey depth and cadence. The parameters that enter the model are the Galactic structure, the distribution of magnetically active and inactive M dwarfs, and the flare frequency distribution (FFD) of both populations. The FFD is a function of spectral type, activity, and Galactic height. Although inactive M dwarfs make up the majority of stars in a magnitude-limited survey, the FFD of inactive stars is very poorly constrained. We have organized a flare monitoring campaign comprising hundreds of hours of new observations from both the ground and space to better constrain flare rates. Incorporating the new observations into our model provides more accurate predictions of stellar variability caused by flares on M dwarfs. We pay particular attention to the likelihood of flares appearing as optical transients (i.e., host star not seen in quiescent data).

Stellar flares are powered by magnetic reconnection events. Although most last for only a few minutes, they have been observed to last for up to 8 hours \citep{Kowalski2010}. Their strongest observational signal is a flux increase in the blue/near-UV, which can be several magnitudes for the biggest flares. Because M dwarfs account for such a large fraction of the stars in the Galaxy, photometric variability on these stars will have an significant signature in time-domain surveys such as Pan-STARRS \citep{Kaiser2004}, PTF \citep{Law2009}, and LSST \citep{LSST2009}. M dwarfs are an excellent place to look for extrasolar planets because planets have relatively large RV effects and transit depths on low-mass stars. Many planets, including several super-Earths, have been discovered orbiting M dwarfs \citep[e.g., ][]{Udry2007,Charbonneau2009,Correia2010}. However, magnetic activity on M dwarfs can hamper planet searches, since starspots and flares introduce both photometric noise and radial velocity jitter \citep{Wright2005,Basri2010,L'opez-Santiago2010}. At the same time, the intense flux of high energy photons and particles caused by flares may have a significant impact on the atmospheres of Earth-like planets in the close-in habitable zones of M dwarfs \citep{Buccino2007,Segura2010}. Our group has used the serendipitous observations of flares in surveys to make estimates of the flare rate. \citet{Kowalski2009} searched for flares in the $\sim$80 photometric epochs from the SDSS Stripe 82 survey \citep{Ivezi'c2007}, while \citet{Hilton2010} identified flares in the sample of M dwarfs with SDSS spectroscopy \citep{West2008}. These studies confirmed that while the flare rate increases with later spectral subtype, the flare energy on earlier type stars was higher. We also found that the flare stars were preferentially closer to the Galactic midplane, which we interpret as an age effect, with younger stars flaring more frequently. Although the occurrence of an individual flare cannot be predicted, dedicated photometric monitoring campaigns of individual stars can yield measurements of the frequency of flares of various energies, also called the flare frequency distribution (FFD). \citet{Moffett1974} and \citet{Lacy1976} obtained FFDs on several of the most well-known and most prolific flare stars in the Solar neighborhood, nearly all of which are active mid-type M dwarfs. We have extended these measurements of the FFD to inactive stars, as well as both early- and late-type M dwarfs.

We have collected nearly 500 hours of flare monitoring observations to make the first measurements of the FFD of inactive early- and mid-type M dwarfs and for active late-type M dwarfs. The FFD is best represented as a linear relation between the log of the quiescent flare energy and the log of the frequency. We find that flares on earlier type stars are more energetic than on later type stars, and that inactive stars flare less frequently than active stars. We incorporate the measured FFDs into our Galactic M dwarf flare model, which can be used to predict the number and magnitude of flares that will be seen in any survey. To make this prediction, we also adopt a flare shape, an active fraction (which is a function of spectral type and distance from the Galactic plane), a luminosity function, and a Galactic model. We draw from the appropriate FFD to generate a light curve for each star along a given sightline. Using the survey parameters from SDSS Stripe 82, our model was able to reproduce both the number and the magnitude distribution of flares found by \citet{Kowalski2009}. Applying our model to the survey parameters of LSST predicts tens of flares $>0.1$ magnitudes and several flares $>1.0$ magnitudes in each {\it u}-band exposure. We can also predict the probability of observing a large flare occurring on a star that is too faint to be seen during quiescence. These large flares, which appear as optical transients in the survey, have a few percent probability of being observed in each {\it u}-band exposure, although the uncertainties are large. Finally, when large numbers of flares are measured in time domain surveys such as Pan-STARRS, PTF, and LSST, our model can be used to precisely determine the FFD of stars in very small bins of spectral type or color and distance from the Galactic plane. This will inform our knowledge of how the magnetic field of M dwarfs changes with spectral type (in particular how it changes across the fully convective boundary), with age, and with activity level.

1012.0577

1012

1012.0154.txt

The recent formulation of the relativistic Thomas-Fermi model within the Feynman-Metropolis-Teller theory for compressed atoms is applied to the study of general relativistic white dwarf equilibrium configurations. The equation of state, which takes into account the $\beta$-equilibrium, the nuclear and the Coulomb interactions between the nuclei and the surrounding electrons, is obtained as a function of the compression by considering each atom constrained in a Wigner-Seitz cell. The contribution of quantum statistics, weak, nuclear, and electromagnetic interactions is obtained by the determination of the chemical potential of the Wigner-Seitz cell. The further contribution of the general relativistic equilibrium of white dwarf matter is expressed by the simple formula $\sqrt{g_{00}}\mu_{\rm ws}=$ constant, which links the chemical potential of the Wigner-Seitz cell $\mu_{\rm ws}$ with the general relativistic gravitational potential $g_{00}$ at each point of the configuration. The configuration outside each Wigner-Seitz cell is strictly neutral and therefore no global electric field is necessary to warranty the equilibrium of the white dwarf. These equations modify the ones used by Chandrasekhar by taking into due account the Coulomb interaction between the nuclei and the electrons as well as inverse $\beta$-decay. They also generalize the work of Salpeter by considering a unified self-consistent approach to the Coulomb interaction in each Wigner-Seitz cell. The consequences on the numerical value of the Chandrasekhar-Landau mass limit as well as on the mass-radius relation of $^4$He, $^{12}$C, $^{16}$O and $^{56}$Fe white dwarfs are presented. All these effects should be taken into account in processes requiring a precision knowledge of the white dwarf parameters.

\label{sec:1} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% The necessity of introducing the Fermi-Dirac statistics in order to overcome some conceptual difficulties in explaining the existence of white dwarfs leading to the concept of degenerate stars was first advanced by R.~H.~Fowler in a classic paper \cite{fowler26}. Following that work, E.~C.~Stoner \cite{stoner29} introduced the effect of special relativity into the Fowler considerations and he discovered the critical mass of white dwarfs \footnote{In doing this, Stoner used what later became known as the exclusion principle, generally attributed in literature to Wolfgang Pauli. For a lucid and scientifically correct historical reconstruction of the contributions to the critical mass concept see \cite{nauenberg08}. For historical details about the exclusion principle see also \cite{heilbron83}.} \begin{equation}\label{eq:maxmassStoner} M^{\rm Stoner}_{\rm crit} = \frac{15}{16} \sqrt{5 \pi} \frac{M_{\rm Pl}^3}{\mu^2 m^2_n} \approx 3.72\frac{M_{\rm Pl}^3}{\mu^2 m^2_n} \, , \end{equation} where $M_{\rm Pl} = \sqrt{\hbar c/G} \approx 10^{-5}$ g is the Planck mass, $m_n$ is the neutron mass, and $\mu = A/Z \approx 2$ is the average molecular weight of matter which shows explicitly the dependence of the critical mass on the chemical composition of the star. Following the Stoner's work, S.~Chandrasekhar \cite{chandrasekhar31} \footnote{At the time a 20 years old graduate student coming to Cambridge from India.} pointed out the relevance of describing white dwarfs by using an approach, initiated by E.~A.~Milne \cite{milne30}, of using the mathematical method of the solutions of the Lane-Emden polytropic equations \cite{emdenbook}. The same idea of using the Lane-Emden equations taking into account the special relativistic effects to the equilibrium of stellar matter for a degenerate system of fermions, came independently to L.~D.~Landau \cite{landau32}. Both the Chandrasekhar and Landau treatments were explicit in pointing out the existence of the critical mass \begin{equation}\label{eq:maxmass} M^{\rm Ch-L}_{\rm crit} = 2.015 \frac{\sqrt{3 \pi}}{2} \frac{M_{\rm Pl}^3}{\mu^2 m^2_n}\approx 3.09\frac{M_{\rm Pl}^3}{\mu^2 m^2_n} \, , \end{equation} where the first numerical factor on the right hand side of Eq.~(\ref{eq:maxmass}) comes from the boundary condition $-(r^2 d u/dr)_{r=R} = 2.015$ (see last entry of Table 7 on Pag.~80 in \cite{emdenbook}) of the $n=3$ Lane-Emden polytropic equation. Namely for $M > M^{\rm Ch-L}_{\rm crit}$, no equilibrium configuration should exist. %This unexpected result created a wave of emotional reactions: Landau rejected the idea of the existence of such a critical mass as a ``ridiculous tendency'' \cite{landau32}. Chandrasekhar was confronted by a lively dispute with A.~Eddington on the basic theoretical assumptions he adopted \footnote{The dispute reached such a heated level that Chandrasekhar was confronted with the option either to change field of research or to leave Cambridge. As is well known he chose the second option transferring to Yerkes Observatory near Chicago where he published his results in his classic book \cite{chandrasekharbook}.} (see \cite{wali82} for details). Some of the basic assumptions adopted by Chandrasekhar and Landau in their idealized approach were not justified e.g. the treatment of the electron as a free-gas without taking into due account the electromagnetic interactions, as well as the stability of the distribution of the nuclei against the gravitational interaction. It is not surprising that such an approach led to the criticisms of Eddington who had no confidence of the physical foundation of the Chandrasekhar work \footnote{It goes to Eddington credit, at the time Plumian Professor at Cambridge, to have allowed the publication of the Chandrasekhar work although preceded by his own critical considerations \cite{eddington35}.}. It was unfortunate that the absence of interest of E.~Fermi on the final evolution of stars did not allow Fermi himself to intervene in this contention and solve definitely these well-posed theoretical problems \cite{ruffinibook}. Indeed, we are showing in this article how the solution of the conceptual problems of the white dwarf models, left open for years, can be duly addressed by considering the relativistic Thomas-Fermi model of the compressed atom (see Subsec.~\ref{subsec:reltf} and Sec.~\ref{sec:4}). The original work on white dwarfs was motivated by astrophysics and found in astrophysics strong observational support. The issue of the equilibrium of the electron gas and the associated component of nuclei, taking into account the electromagnetic, the gravitational and the weak interactions is a theoretical physics problem, not yet formulated in a correct special and general relativistic context. One of the earliest alternative approaches to the Chandrasekhar-Landau work was proposed by E.~E.~Salpeter in 1961 \cite{salpeter61}. He followed an idea originally proposed by Y.~I.~Frenkel \cite{frenkel28}: to adopt in the study of white dwarfs the concept of a Wigner-Seitz cell. Salpeter introduced to the lattice model of a point-like nucleus surrounded by a uniform cloud of electrons, corrections due to the non-uniformity of the electron distribution (see Subsec.~\ref{subsec:salpeter} for details). In this way Salpeter \cite{salpeter61} obtained an analytic formula for the total energy in a Wigner-Seitz cell and derived the corresponding equation of state of matter composed by such cells, pointing out explicitly the relevance of the Coulomb interaction. The consequences of the Coulomb interactions in the determination of the mass and radius of white dwarfs, was studied in a subsequent paper by T.~Hamada and E.~E.~Salpeter \cite{hamada61} by using the equation of state constructed in \cite{salpeter61}. They found that the critical mass of white dwarfs depends in a nontrivial way on the specific nuclear composition: the critical mass of Chandrasekhar-Landau which depends only on the mass to charge ratio of nuclei $A/Z$, now depends also on the proton number $Z$. This fact can be seen from the approximate expression for the critical mass of white dwarfs obtained by Hamada and Salpeter \cite{hamada61} in the ultrarelativistic limit for the electrons \begin{equation}\label{eq:maxmassHS} M^{\rm H\&S}_{\rm crit} = 2.015 \frac{\sqrt{3 \pi}}{2} \frac{1}{\mu^2_{\rm eff}} \frac{M_{\rm Pl}^3}{m^2_n}\, , \end{equation} where \begin{equation}\label{eq:mueff} \mu_{\rm eff} = \mu \left(\frac{P_{\rm S}}{P_{\rm Ch}}\right)^{-3/4}\, , \end{equation} being $P_{\rm S}$ the pressure of the Wigner-Seitz cell obtained by Salpeter in \cite{salpeter61} (see Subsec.~\ref{subsec:salpeter}) and $P_{\rm Ch}$ is the pressure of a free-electron fluid used by Chandrasekhar (see Subsec.~\ref{subsec:uniform}). The ratio $P_{\rm S}/P_{\rm Ch}$ is a function of the number of protons $Z$ (see Eq.~(20) in \cite{salpeter61}) and it satisfies $P_{\rm S}/P_{\rm Ch} < 1$. Consequently, the effective molecular weight satisfies $\mu_{\rm eff} > \mu$ and the critical mass of white dwarfs turns to be smaller than the original one obtained by Chandrasekhar-Landau (see Eq.~(\ref{eq:maxmass})). In the mean time, the problem of the equilibrium gas in a white dwarf taking into account possible global electromagnetic interactions between the nucleus and the electrons was addressed by E.~Olson and M.~Bailyn in \cite{olson75,olson76}. They well summarized the status of the problem: ``\emph{Traditional models for the white dwarf are non-relativistic and electrically neutral ... although an electric field is needed to support the pressureless nuclei against gravitational collapse, the star is treated essentially in terms of only one charge component, where charge neutrality is assumed }''. Their solution to the problem invokes the breakdown of the local charge neutrality and the presence of an overall electric field as a consequence of treating also the nuclei inside the white dwarf as a fluid. They treated the white dwarf matter through a two-fluid model not enforcing local charge neutrality. The closure equation for the Einstein-Maxwell system of equations was there obtained from a minimization procedure of the mass-energy of the configuration. This work was the first pointing out the relevance of the Einstein-Maxwell equations in the description of an astrophysical system by requiring global and non local charge neutrality. As we will show here, this interesting approach does not apply to the case of white dwarfs. It represents, however, a new development in the study of neutron stars (see e.g.~\cite{PLB2011}) An alternative approach to the Salpeter treatment of a compressed atom was reconsidered in \cite{gursky2000} by applying for the first time to white dwarfs a relativistic Thomas-Fermi treatment of the compressed atom introducing a finite size nucleus within a phenomenological description (see also \cite{bertone2000}). Recently, the study of a compressed atom has been revisited in \cite{2011PhRvC..83d5805R} by extending the global approach of Feynman, Metropolis and Teller \cite{feynman49} taking into account weak interactions. This treatment takes also into account all the Coulomb contributions duly expressed relativistically without the need of any piecewise description. The relativistic Thomas-Fermi model has been solved by imposing in addition to the electromagnetic interaction also the weak equilibrium between neutrons, protons and electrons self-consistently. This presents some conceptual differences with respect to previous approaches and can be used in order both to validate and to establish their limitations. In this article we apply the considerations presented in \cite{2011PhRvC..83d5805R} of a compressed atom in a Wigner-Seitz cell to the description of non-rotating white dwarfs in general relativity. This approach improves all previous treatments in the following aspects: \begin{enumerate} %% \item In order to warranty self-consistency with a relativistic treatment of the electrons, the point-like assumption of the nucleus is abandoned introducing a finite sized nucleus \cite{2011PhRvC..83d5805R}. We assume for the mass as well as for charge to mass ratio of the nucleus their experimental values instead of using phenomenological descriptions based on the semi-empirical mass-formula of Weizsacker (see e.g.~\cite{gursky2000,bertone2000}). %% \item The electron-electron and electron-nucleus Cou\-lomb interaction energy is calculated without any approximation by solving numerically the relativistic Thomas-Fermi equation for selected energy-densities of the system and for each given nuclear composition. %% \item The energy-density of the system is calculated taking into account the contributions of the nuclei, of the Coulomb interactions as well as of the relativistic electrons; the latter being neglected in all previous treatments. This particular contribution turns to be very important at high-densities and in particular for light nuclear compositions e.g.~$^4$He and $^{12}$C. %% \item The $\beta$-equilibrium between neutrons, protons, and electrons is also taken into account leading to a self-consistent calculation of the threshold density for triggering the inverse $\beta$-decay of a given nucleus. %% \item The structure of the white dwarf configurations is obtained by integrating the general relativity equations of equilibrium. %% \item Due to 4) and 5) we are able to determine if the instability point leading to a maximum stable mass of the non-rotating white dwarf is induced by the inverse $\beta$-decay instability of the composing nuclei or by general relativistic effects. %% \end{enumerate} Paradoxically, after all this procedure which takes into account many additional theoretical features generalizing the Chandrasekhar-Landau and the Hamada and Salpeter works, a most simple equation is found to be fulfilled by the equilibrium configuration in a spherically symmetric metric. Assuming the metric \begin{equation}\label{eq:metric} ds^2 = e^{\nu(r)} c^2 dt^2 - e^{\lambda(r)}dr^2 - r^2 d\theta^2 - r^2 \sin^2 \theta d\varphi^2\, , \end{equation} we demonstrate how the entire system of equations describing the equilibrium of white dwarfs, taking into account the weak, the electromagnetic and the gravitational interactions as well as quantum statistics all expressed consistently in a general relativistic approach, is simply given by \begin{equation}\label{eq:conslaw} \sqrt{g_{00}} \mu_{\rm ws} = e^{\nu(r)/2}\mu_{\rm ws}(r) = {\rm constant}\, , \end{equation} which links the chemical potential of the Wigner-Seitz cell $\mu_{\rm ws}$, duly solved by considering the relativistic Feynman-Metropolis-Teller model following \cite{2011PhRvC..83d5805R}, to the general relativistic gravitational potential at each point of the configuration. The overall system outside each Wigner-Seitz cell is strictly neutral and no global electric field exists, contrary to the results reported in \cite{olson76}. The same procedure will apply as well to the case of neutron star crusts. The article is organized as follows. In Sec.~\ref{sec:2} we summarize the most common approaches used for the description of white dwarfs and neutron star crusts: the uniform approximation for the electron fluid (see e.g.~\cite{chandrasekhar31}); the often called lattice model assuming a point-like nucleus surrounded by a uniform electron cloud (see e.g.~\cite{baym71a}); the generalization of the lattice model due to Salpeter \cite{salpeter61}; the Feynman, Metropolis and Teller approach \cite{feynman49} based on the the non-relativistic Thomas-Fermi model of compressed atoms and, the relativistic generalization of the Feynman-Metropolis-Teller treatment recently formulated in \cite{2011PhRvC..83d5805R}. In Sec.~\ref{sec:3} we formulate the general relativistic equations of equilibrium of the system and show how, from the self-consistent definition of chemical potential of the Wigner-Seitz cell and the Einstein equations, comes the equilibrium condition given by Eq.~(\ref{eq:conslaw}). In addition, we obtain the Newtonian and the first-order post-Newtonian equations of equilibrium. Finally, we show in Sec.~\ref{sec:4} the new results of the numerical integration of the general relativistic equations of equilibrium and discuss the corrections to the Stoner critical mass $M^{\rm Stoner}_{\rm crit}$, to the Chandrasekhar-Landau mass limit $M^{\rm Ch-L}_{\rm crit}$, as well as to the one of Hamada and Salpeter $M^{\rm H\&S}_{\rm crit}$, obtained when all interactions are fully taken into account through the relativistic Feynman-Metropolis-Teller equation of state \cite{2011PhRvC..83d5805R}. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%

\label{sec:5} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% We have addressed the theoretical physics aspects of the white dwarf configurations of equilibrium, quite apart from the astrophysical application. %In the introduction we have recalled how the study of white dwarfs has often stimulated and taken advantage of crucial progress in theoretical physics and applied mathematics. It is clear that the early considerations of the critical mass of a white dwarf were routed in the concept of quantum statistics and the fermion exclusion principle considered by Fowler \cite{fowler26} and Stoner \cite{stoner24} leading to the critical mass (\ref{eq:maxmassStoner}) \cite{stoner29}. % %The following progress was made by Chandrasekhar \cite{chandrasekhar31} and Landau \cite{landau32} adopting the work by Emden \cite{emdenbook} on the solution of the nonlinear Lane-Emden polytropic differential equations. They obtained the critical mass given by Eq.~(\ref{eq:maxmass}). % %It was Salpeter \cite{salpeter61} and later Hamada and Salpeter \cite{hamada61} who brought to full fruition the additional conceptual theoretical physics concept of Wigner-Seitz cells. Salpeter adopted the Wigner-Seitz cell for the description of white dwarf material and studied the perturbation to the uniform electron distribution, given by the Coulomb interactions and their special relativity corrections. The value of the critical mass, although obtained only through numerical integration, can be expressed approximately by Eq.~(\ref{eq:maxmassHS}). %Still many inconsistencies existed in the theoretical model. The recently accomplished description of a compressed atom within the global approach of the relativistic Feynman, Metropolis and Teller \cite{2011PhRvC..83d5805R} has been here solved within the Wigner-Seitz cell and applied to the construction of white dwarfs in the framework of general relativity. From a theoretical physics point of view, this is the first unified approach of white dwarfs taking into account consistently the gravitational, the weak, the strong and the electromagnetic interactions, and it answers open theoretical physics issues in this matter. No analytic formula for the critical mass of white dwarfs can be derived and, on the contrary, the critical mass can obtained only through the numerical integration of the general relativistic equations of equilibrium together with the relativistic Feynman-Metropolis-Teller equation of state. The value of the critical mass and the radius of white dwarfs in our treatment and in the Hamada and Salpeter \cite{hamada61} treatment becomes a function of the composition of the star. Specific examples have been given in the case of white dwarfs composed of $^4$He, $^{12}$C, $^{16}$O and $^{56}$Fe. The results of Chandrasekhar, of Hamada and Salpeter and ours have been compared and contrasted (see Table \ref{tab:mcrit} and Figs.~\ref{fig:MrhoHe}--\ref{fig:MRFe}). The critical mass is a decreasing function of $Z$ and Coulomb effects are more important for heavy nuclear compositions. The validity of the Salpeter approximate formulas increases also with $Z$, namely for heavy nuclear compositions the numerical values of the masses as well as of the radii of white dwarfs obtained using the Salpeter equation of state are closer to the ones obtained from the full numerical integration of the general relativistic treatment presented here. Turning now to astrophysics, the critical mass of white dwarfs is today acquiring a renewed interest in view of its central role in the explanation of the supernova phenomena \cite{phillips93,riess98,perlmutter99,riess04}. The central role of the critical mass of white dwarfs as related to supernova was presented by F.~Hoyle and W.~A.~Fowler \cite{hoyle60} explaining the difference between type I and type II Supernova. This field has developed in the intervening years to a topic of high precision research in astrophysics and, very likely, both the relativistic and the Coulomb effects outlined in this article will become topic of active confrontation between theory and observation. For instance, the underestimate of the mass and the radius of low density white dwarfs within the Hamada and Salpeter treatment \cite{hamada61} (see Figs.~\ref{fig:MrhoHe}--\ref{fig:MRFe}) leads to the possibility of a direct confrontation with observations in the case of low mass white dwarfs e.g. the companion of the Pulsar J1141-6545 \cite{kramerprivate}. %Paradoxically, the concept of critical mass was not pursued by Chandrasekhar in order to explain the supernova phenomena, on the contrary, Chandrasekhar purported on the role of the critical mass in discriminating star masses leading to the formation of a white dwarf versus stars never reaching a configuration of equilibrium due to radiation pressure \cite{arnettprivate,giacconi78} \footnote{Chandrasekhar, in an interview with S.~Weart \cite{weart77}, recognized ``... at first I didn’t understand what this limit meant and I didn’t know how it would end, and how it related to the 3/2 low mass polytropes ...}. We have finally obtained a general formula in Eq.~(\ref{eq:conslaw2}) as a ``first integral'' of the general relativistic equations of equilibrium. This formula relates the chemical potential of the Wigner-Seitz cells, duly obtained from the relativistic Feynman-Metropolis-Teller model \cite{2011PhRvC..83d5805R} taking into account weak, nuclear and electromagnetic interactions, to the general relativistic gravitational potential at each point of the configuration. Besides its esthetic value, this is an important tool to examine the radial dependence of the white dwarf properties and it can be also applied to the crust of a neutron star as it approaches to the physical important regime of neutron star cores. The formalism we have introduced allows in principle to evaluate subtle effects of a nuclear density distribution as a function of the radius and of the Fermi energy of the electrons and of the varying depth of the general relativistic gravitational potential. The theoretical base presented in this article establishes also the correct framework for the formulation of the more general case when finite temperatures and magnetic fields are present. This treatment naturally opens the way to a more precise description of the crust of neutron stars, which will certainly become an active topic of research in view of the recent results by S. Goriely et al. \cite{2011A&A...531A..78G,2011arXiv1107.0899G} and by J.~M.~Pearson et al. \cite{2011PhRvC..83f5810P} on the importance of the Coulomb effects in the r-process nucleosynthesis of the crust material during its post-ejection evolution in the process of gravitational collapse and/or in the merging of neutron star binaries. %We have explicitly discussed many theoretical issues open for years on white dwarfs, including the legitimately posed by A.~Eddington \cite{eddington35}.

1012.0154

1012

1012.4929_arXiv.txt

Delayed detonations of Chandrasekhar-mass white dwarfs (WDs) have been very successful in explaining the spectra, light curves, and the width-luminosity relation of spectroscopically normal Type Ia supernovae (SNe~Ia). The ignition of the thermonuclear deflagration flame at the end of the convective carbon ``simmering'' phase in the core of the WD is still not well understood and much about the ignition kernel distribution remains unknown. Furthermore, the central density at the time of ignition depends on the still uncertain screened carbon fusion reaction rates, the accretion history and cooling time of the progenitor, and the composition. We present the results of twelve high-resolution three-dimensional delayed detonation SN~Ia explosion simulations that employ a new criterion to trigger the deflagration to detonation transition (DDT). The simulations fall into into three ignition categories: relatively bright SNe with 5 ignition kernels and a weak deflagration phase (three different central densities), relatively dim SNe with 1600 ignition kernels and a strong deflagration phase (three different central densities) and intermediate SNe with 200 ignition kernels (six different central densities). All simulations trigger our DDT criterion and the resulting delayed detonations unbind the star. We find a trend of increasing iron group element (IGE) production with increasing central density for all three categories. The total \nuc{56}{Ni} yield, however, remains more or less constant, even though increased electron captures at high density result in a decreasing \nuc{56}{Ni} mass fraction of the IGE material. We attribute this to an approximate balance of \nuc{56}{Ni} producing and destroying effects. The deflagrations that were ignited at higher density initially have a faster growth rate of subgrid-scale turbulence. Hence, the effective flame speed increases faster, which triggers the DDT criterion earlier, at a time when the central density of the expanded star is higher. This leads to an overall increase of IGE production, which off-sets the percental reduction of \nuc{56}{Ni} due to neutronization.

\label{sec:int} SNe Ia have come to fame as the Universe's most luminous standardizable candles -- crucial ingredients to the study of dark energy and cosmology \citep[e.g.][]{riess1998a,schmidt1998a}. A limiting factor on the precision of using SNe~Ia as distance indicators is the inherent scatter in their normalized light curves \citep[e.g.][]{wood-vasey2007a}. A better understanding of the intrinsic variation of supernova brightnesses and spectra is needed \citep[e.g.][]{albrecht2006a,miknaitis2007a}. Simulations of SN Ia explosions are already being used to aid in improving the precision of cosmological distance measurements based on supernovae in the future \citep[e.g.][]{blondin2011a}. In addition, SNe~Ia also play a critical role in galaxy gas kinematics \citep[e.g.][]{scannapieco2008a}, positron production \citep[e.g.][]{chan1993a}, and chemical evolution \citep[e.g.][]{matteucci1986a}. Detailed modeling of the explosions is therefore useful for understanding the origin of the Galactic 511 keV line, the origin and evolution of heavy elements, and kinetic supernova feedback, and measuring the Hubble parameter as a function of redshift. The standard model of SNe Ia relies on the nuclear fusion of the initial composition (predominantly \nuc{12}{C} and \nuc{16}{O}) of a massive white dwarf (WD) star to more tightly bound nuclei to power the explosion \citep{hoyle1960a}. The exact nature of the progenitor systems and details of the dynamics of the nuclear burning processes however are not known. Among the leading scenarios are the Chandrasekhar-mass models, in which a WD accretes matter from a companion star and grows in mass to near the Chandrasekhar limit until pycnonuclear carbon fusion reactions \citep{cameron1959a} start taking place. Once carbon fusion reactions produce more energy than is carried away by neutrino losses, the core becomes convective and when the nuclear burning time of a fluid element becomes shorter than the eddy turnover time a deflagration flame may be born \citep[e.g.][]{woosley1990a}. Numerical simulations of the convective stage leading up to the ignition of the deflagration were performed by \citet{hoeflich2002a}, \citet{kuhlen2006a}, \citet{piro2008b}, \citet{piro2008c}, and \citet{zingale2009a}. The central density of the WD decreases significantly during the simmering phase between the onset of carbon burning and the ignition of the deflagration \citep[e.g.][]{lesaffre2006a,piro2008c}. Those calculations, however, are not taking electron captures and the URCA process correctly into account, and some uncertainty in the evolution remains. The rate of the screened \nuc{12}{C}--\nuc{12}{C} fusion reaction is still quite uncertain \citep[e.g.][]{itoh2003a,jiang2007a,gasques2005a,gasques2007a}. The central density at the time of ignition, however, depends only mildly on the exact value of this reaction rate \citep{cooper2009a, iapichino2010a}. More important is the initial mass and the accretion and cooling history of the WD, which determines the thermodynamic state of the interior. This results in a range of possible central densities at ignition, from less than $2 \times 10^9$ to over $5 \times 10^9 \gcc$ \citep{lesaffre2006a}. Metallicity has a considerable impact on the supernova brightness \citep[e.g.][]{timmes2003a,travaglio2005a,bravo2010a}. In contrast, the ignition density has been shown to depend rather weakly on metallicity and the CO ratio \citep{lesaffre2006a}. If the initial deflagration flame can transition into a detonation \citep[e.g.][]{khokhlov1997a,roepke2007d,woosley2007a,woosley2009a}, then good agreement of the models with observations can be obtained \citep[e.g.][]{roepke2007b,bravo2008a,kasen2009a}. A successful explosion model has to reproduce the observed range of peak absolute magnitudes (i.e. \nuc{56}{Ni} masses) and the width-luminosity relation and scatter thereabout. Furthermore, the observed correlation between the brightness of an event and the delay time or age of the host stellar population has to be explained \citep[e.g.][]{gallagher2008a}. Recently, a connection between age of the host stellar population and SN~Ia brightness was proposed via the effect of longer cooling times on the ignition density \citep{krueger2010a}. Varying the central density for 150 two dimensional delayed detonation supernova simulations within the statistical ignition framework presented in \citet{townsley2009a}, the authors found that the \nuc{56}{Ni} yield decreased with increasing central density, while the total iron group element (IGE) yield remains roughly constant. This is attributed to increased production of stable isotopes (such as e.g. \nuc{54}{Fe} or \nuc{58}{Ni}) due to increased neutronization via electron captures at the higher densities. There are, however, at least three competing effects that influence the \nuc{56}{Ni} mass produced in a delayed detonation SN. \begin{enumerate} \item Electron capture rates on protons and iron-group isotopes under electron degenerate condisitons are strongly increasing with density \citep[e.g.][]{langanke2001a}. Consequently, a distribution of nuclei in nuclear statistical equilibrium at high density neutronizes at a much faster rate than one at lower density \citep[e.g.][]{seitenzahl2009a}, which acts to lower the \nuc{56}{Ni} mass. \item Near Chandrasekhar-mass WDs in hydrostatic equilibrium with a higher central density are more compact, i.e. significantly smaller and slightly more massive and tighly bound. This may translate into a more compact WD at the time of the first DDT, which could lead to an overall larger part of the WD being burned to IGEs, which acts to raise the \nuc{56}{Ni} mass. \item Deflagrations evolve differently at higher gravitational acceleration $g$ \citep{khokhlov1995a,zhang2007a}. From linear stability analysis, the Rayleigh-Taylor temporal growth rate scales with $\sqrt{g}$. The different flame evolution and turbulence generation could have an effect on the DDT (e.g. the transition density), which, depending on the different degree of ``pre-expansion'', could either lower or raise the \nuc{56}{Ni} mass. \end{enumerate} The effect of variations in the central density of the WD on \emph{pure deflagrations} has been explored in three-dimensional models before \citep{roepke2006b}. Here, we present the results of twelve high-resolution three-dimensional \emph{delayed detonation} SN~Ia simulations (that employ a new DDT criterion, see Section \ref{sec:ddt}) for three different ignition configurations and a range of central densities. We find that, for the same spatial ignition spark distributions, the \nuc{56}{Ni} yield remains more or less constant a function of central density at ignition. The deflagrations that were ignited at higher density produce subgrid-scale turbulence at a higher rate, which triggers the DDT criterion earlier when the central density of the star is higher. This leads to an overall increase of IGE production as well as enhanced electron captures. Even though the mass in \nuc{56}{Ni} comprises a smaller fraction of the mass that has burned to IGEs, the overall \nuc{56}{Ni} yield remains roughly constant since more total mass in IGEs is produced in the detonation. Only the cases where much of the IGEs are produced in the deflagration phase show a trend of decreasing \nuc{56}{Ni} with central density. In Section~\ref{sec:sims} we introduce our setup and briefly review the computational methods, in Section~\ref{sec:results} and \ref{sec:discussion} we present and discuss the results, and in Section~\ref{sec:conclusions} we conclude.

\label{sec:conclusions} We have performed twelve three-dimensional hydrodynamical simulations for delayed detonation SNe~Ia for a range of central densities and ignition conditions. We find a trend of increasing IGE production with central density within each set of ignition conditions. This is because the high central density WDs are more compact and the flame evolves faster; the DDT occurs sooner when more unburned material is still above the density threshold ($\approx 10^7 \gcc$) where a detonation will still produce IGE. In spite of the larger IGE mass, the more vigorous neutronization occurring in the high density models during the deflagration phase yields \nuc{56}{Ni} massest hat are more or less constant with $\rho_{\mathrm{c}}$ for the brighter SNe. Only dim SNe, which have a strong deflagration phase and expansion prior to the DDT, exhibit a trend of decreasing \nuc{56}{Ni} mass with increasing density, since the increased neutronization in the deflagration phase cannot be compensated for by the relatively weak detonation phase. This trend, however, is of secondary importance when compared to the effects of varying the ignition kernel distribution. For a given ignition kernel spatial distribution, the central density therefore influences the brightness of the supernova event only as a secondary parameter. From the works of \citet{townsley2009a} and \citet{bravo2010a}, it appears that the same holds for composition, i.e. metallicity and C/O ratio. Indeed, based on an analysis of high-quality V and B-band light curves of SNe~Ia from the Carnegie Supernova Project, \citet{hoeflich2010a} propose that the composition and central density are two independent secondary parameters for SN~Ia light curves. In light of the importance of the ignition configuration of the deflagration for the brightness of the SN, it is most crucial to establish how the central density at ignition (cooling time) and metallicity affect the statistical properties (notably number and location) of the ignition sparks themselves, and not their respective direct effects on the outcome of an explosion once a random ignition spark distribution was chosen. One should therefore aim to quantify which effect composition, cooling and accretion history have on the ignition process, for example by mapping them into the exponentiation parameter $C_{\mathrm{e}}$ of the stochastic ignition prescription of \citet{schmidt2006a}. This would require a better understanding of the physics leading up to ignition, including the nature of the convection and effects of electron captures and the convective URCA process.

1012.4929

1012

1012.2022_arXiv.txt

In this work we study the constraints on the dark matter interaction with the standard model particles, from the observations of dark matter relic density, the direct detection experiments of CDMS and XENON, and the indirect detection of the $\bar{p}/p$ ratio by PAMELA. A model independent way is adopted in the study by constructing the effective interaction operators between dark matter and standard model particles. The most general 4-fermion operators are investigated. We find that the constraints from different observations are complementary with each other. Especially the spin independent scattering gives very strong constraints for corresponding operators. In some cases the indirect detection of $\bar{p}/p$ data can actually be more sensitive than the direct detection or relic density for light dark matter ($\lesssim 70$ GeV).

The existence of a significant component of nonbaryonic dark matter (DM) in the Universe has been well confirmed by astrophysical observations \cite{Tegmark:2006az,Komatsu:2008hk,Komatsu:2010fb} in recent years, however the nature of this substance remains unclear. Since there is no candidate for DM in the Standard Model (SM) of particle physics, it implies the existence of new physics beyond the SM. Probe of the microscopic identity and properties of DM has become one of the key problems in particle physics and cosmology (for reviews of DM, see, for instance, \cite{kolb-turner,Jungman:1995df,Bertone:2004pz,Murayama:2007ek,Feng:2010gw}). Among a large amount of theoretical models, a well-motivated candidate for DM is the weakly interacting massive particle (WIMP). This WIMP must be stable, nonrelativistic, electrically neutral, and colorless. If the mass of WIMP is from a few GeV to TeV and the interaction strength is of the weak scale, they can naturally yield the observed relic density of DM, which is often referred to as the WIMP miracle \cite{Feng:2010gw}. A huge variety of new physics models trying to solve the problems of the SM at the weak scale can naturally contain WIMP candidates, such as the supersymmetric models \cite{Jungman:1995df,Goldberg:1983nd,Ellis:1983ew,Kane:1993td}, extra dimensional models \cite{Kolb:1983fm,Cheng:2002ej,Hooper:2007qk,Servant:2002aq,Servant:2002hb,Agashe:2004ci, Agashe:2004bm,Agashe:2007jb}, little Higgs models \cite{Cheng:2004yc,Low:2004xc,Birkedal:2006fz,Freitas:2009jq,Kim:2009dr}, left-right symmetric models \cite{Dolle:2007ce,Guo:2008hy,Guo:2008si,Guo:2010vy}, and many other theoretical scenarios. The above mentioned specific models are well-motivated, however they still lack experimental support. We do not know whether nature really behaves like one of them or some other yet unconsidered theories. Moreover, in case the DM particle is the only new particle within the reach of LHC and other new particle species are much heavier than DM, it will be very difficult to tell which model the DM particle belongs to. Additionally, it is possible that the DM may be first observed by direct or indirect detection experiments. These early observations may only provide information about some general properties of the DM particle, and may not be able to distinguish the underlying theories. Therefore, the model-independent studies of the DM phenomenology are particularly important for they may avoid theoretical bias \cite{Birkedal:2004xn,Giuliani:2004uk,Kurylov:2003ra,Beltran:2008xg,Cirelli:2008pk,Shepherd:2009sa}. Recently there have been quite a few papers following such consideration and adopting a model-independent way to study various phenomenologies related with DM \cite{Cao:2009uv,Cao:2009uw,Beltran:2010ww,Fitzpatrick:2010em,Goodman:2010yf,Bai:2010hh, Goodman:2010ku,Goodman:2010qn,Bell:2010ei}. Especially the relic density measured by WMAP \cite{Komatsu:2010fb}, direct detection from CDMS \cite{Ahmed:2009zw}, XENON \cite{Aprile:2010um} and possible collider signals from LHC are considered in these studies. In this work, we first construct the general effective 4-fermion interaction operators between DM particles and the SM particles, which extend the effective fermionic WIMP interactions given in Ref. \cite{Beltran:2008xg}. Here we focus on Dirac fermionic DM. Discussions on scalar and vector DM will be presented in companion papers. We then give updated constraints from the DM relic density within the 7-year WMAP data \cite{Komatsu:2010fb} and the spin-independent WIMP-nucleus elastic scattering searches by CDMS II \cite{Ahmed:2009zw} and XENON100 \cite{Aprile:2010um}, and compare our results with those in Ref. \cite{Beltran:2008xg}. In addition, we present new phenomenological constraints on these effective models from the spin-dependent WIMP-nucleus elastic scattering searches by CDMS \cite{Akerib:2005za} and XENON \cite{Angle:2008we} and the cosmic-ray antiproton-to-proton ratio by PAMELA \cite{Adriani:2010rc}. We find that the constraints from different kinds of experiments are rather comparable. Combination of these constraints provides more information of the effective models. This paper is organized as follows. In Sec. \ref{sec-model} we briefly describe the effective DM models of various 4-fermion interaction operators. In Sec. \ref{sec-relic}, \ref{sec-direct} and \ref{sec-indirect} we explore the constraints on these models from the DM relic density, direct and indirect detection searches, respectively. In Sec. \ref{sec-combine} we discuss the validity region of effective theory and present the combined constraints on the effective coupling constants of these models. Sec. \ref{sec-con} is the conclusion.

} \begin{table}[!htbp] \begin{center} \belowcaptionskip=0.2cm \caption{A summary for Dirac fermionic WIMPs with various effective interactions. The excluded regions of $M_\chi$ given by direct and indirect experiments are indicated.} \label{tab:dirac_sum} \renewcommand{\arraystretch}{1.3} \small \begin{tabular}{ccc} \hline \hline \multicolumn{3}{c}{Universal coupling} \\ Interaction & Direct detection & PAMELA $\bar p / p$ \\ \hline Scalar & Excluded $M_\chi \simeq 10~\mathrm{GeV} - \mathrm{above}~1~\mathrm{TeV}$ & Not sensitive \\ Pseudoscalar & Not sensitive & Excluded $M_\chi \simeq 10 - 70~\mathrm{GeV}$ \\ Vector & Excluded $M_\chi \simeq 10~\mathrm{GeV} - 1~\mathrm{TeV}$ & Excluded $M_\chi \simeq 10 - 70~\mathrm{GeV}$ \\ Axialvector & Not sensitive & Excluded $M_\chi \simeq 10 - 14~\mathrm{GeV}$ \\ Tensor & Not sensitive & Excluded $M_\chi \simeq 10 - 70~\mathrm{GeV}$ \\ Scalar-pseudoscalar & Not sensitive & Not sensitive \\ Pseudoscalar-scalar & Not sensitive & Excluded $M_\chi \simeq 10 - 70~\mathrm{GeV}$ \\ Vector-axialvector & Not sensitive & Excluded $M_\chi \simeq 10 - 70~\mathrm{GeV}$ \\ Axialvector-vector & Not sensitive & Not sensitive \\ Alternative tensor & Not sensitive & Excluded $M_\chi \simeq 10 - 70~\mathrm{GeV}$ \\ Chiral & Excluded $M_\chi \simeq 10~\mathrm{GeV} - 1~\mathrm{TeV}$ & Excluded $M_\chi \simeq 10 - 70~\mathrm{GeV}$ \\ \hline \hline \multicolumn{3}{c}{$G_f \propto m_f$} \\ Interaction & Direct detection & PAMELA $\bar p / p$ \\ \hline Scalar & Excluded $M_\chi \simeq 10 - 185~\mathrm{GeV}$ & Not sensitive \\ Pseudoscalar & Not sensitive & Excluded $M_\chi \simeq 10 - 70~\mathrm{GeV}$ \\ Vector & Not sensitive & Excluded $M_\chi \simeq 10 - 70~\mathrm{GeV}$ \\ Axialvector & Not sensitive & Excluded $M_\chi \simeq 10 - 25~\mathrm{GeV}$ \\ Tensor & Not sensitive & Excluded $M_\chi \simeq 10 - 70~\mathrm{GeV}$ \\ Scalar-pseudoscalar & Not sensitive & Not sensitive \\ Pseudoscalar-scalar & Not sensitive & Excluded $M_\chi \simeq 10 - 70~\mathrm{GeV}$ \\ Vector-axialvector & Not sensitive & Excluded $M_\chi \simeq 10 - 70~\mathrm{GeV}$ \\ Axialvector-vector & Not sensitive & Not sensitive \\ Alternative tensor & Not sensitive & Excluded $M_\chi \simeq 10 - 70~\mathrm{GeV}$ \\ Chiral & Not sensitive & Excluded $M_\chi \simeq 10 - 68~\mathrm{GeV}$ \\ \hline \hline \end{tabular} \end{center} \end{table} In this work we give a general analysis of the 4-fermion interaction between the DM and the standard model particles. We have considered the most general form of the 4-fermion operators and corrected some errors in the previous works. We find that the constraints from DM relic density, DM direct detection and indirect detection of $\bar{p}/p$ data are complementary to each other. Generally, the SI constraints are the most stringent while the SD constraints are quite weak. For light DM ($\lesssim 70$ GeV) the $\bar{p}/p$ data give very strong constraints on the interaction. Assuming that one operator dominates the effective interaction between DM and the SM fermions, we find that some cases get so strong constraints that the universe will be overclosed by DM thermal production. In such cases the DM models are actually excluded assuming a standard cosmology. As a summary, in Tab.~\ref{tab:dirac_sum}, we indicate the excluded regions of $M_\chi$ given by direct and indirect experiments for Dirac fermionic WIMPs with various effective interactions. We find that recent direct detection experiments only exclude some regions of $M_\chi$ for the scalar, vector and chiral interactions with universal couplings, and for the scalar interaction with $G_f \propto m_f$. The PAMELA $\bar p / p$ spectrum, however, excludes some small $M_\chi$ regions ($\lesssim 70~\mathrm{GeV}$) for most of the effective interactions.

1012.2022

1012

1012.5434_arXiv.txt

Radio polarimetry is a three-dimensional statistical problem. The three-dimensional aspect of the problem arises from the Stokes parameters Q, U, and V, which completely describe the polarization of electromagnetic radiation and conceptually define the orientation of a polarization vector in the Poincar\'e sphere. The statistical aspect of the problem arises from the random fluctuations in the source-intrinsic polarization and the instrumental noise. A simple model for the polarization of pulsar radio emission has been used to derive the three-dimensional statistics of radio polarimetry. The model is based upon the proposition that the observed polarization is due to the incoherent superposition of two, highly polarized, orthogonal modes. The directional statistics derived from the model follow the Bingham-Mardia and Fisher family of distributions. The model assumptions are supported by the qualitative agreement between the statistics derived from it and those measured with polarization observations of the individual pulses from pulsars. The orthogonal modes are thought to be the natural modes of radio wave propagation in the pulsar magnetosphere. The intensities of the modes become statistically independent when generalized Faraday rotation (GFR) in the magnetosphere causes the difference in their phases to be large. A stochastic version of GFR occurs when fluctuations in the phase difference are also large, and may be responsible for the more complicated polarization patterns observed in pulsar radio emission.

The four fundamental measurements made in astronomy are the intensity, flux density, or surface brightness of the electromagnetic radiation emitted by a celestial object, the wavelength, or frequency, of the radiation, its location on the sky, and the polarization of the radiation. Measurements of the latter two, location and polarization, follow the statistics of direction. The association of directional statistics with the measurement of location is obvious, but the application of directional statistics to polarization measurements is not immediately apparent until one recalls that the Stokes parameters Q, U, and V describe the orientation of a polarization vector within the Poincar\'e sphere. The Stokes parameter V defines the circular polarization of the radiation and establishes the \lq\lq z-coordinate" of the polarization vector in the Poincar\'e sphere. The Stokes parameters Q and U describe the radiation's linear polarization and establish the vector's x- and y-coordinates, respectively. Here, polarization measurements are shown to follow directional statistics, and these statistics are applied to polarization observations of radio pulsars. Pulsars are rapidly rotating, highly magnetized neutron stars. Their rotation periods range between about 1ms and 10s, and the strength of the magnetic field at their surfaces ranges from $10^8$ G for the oldest pulsars to over $10^{12}$ G for the youngest. A beam of radio emission is emitted from each of the star's magnetic poles. A pulse of radio emission is observed as the star's rotation causes the beam to sweep across an observer's line of sight. Pulsar radio emission is generally thought to originate from charged particles streaming along open magnetic fields lines above the star's magnetic pole, but unlike other astrophysical radiative processes (e.g. synchrotron radiation, maser emission, and thermal radiation), it is poorly understood. Polarization observations of the individiual pulses from pulsars are made in an attempt to understand the radio emission mechanism and to study the propagation of radio waves in ultra-strong magnetic fields. Polarization observations of individual pulses (Lyne et al. 1971; Manchester et al. 1975; Backer \& Rankin 1980; Stinebring et al. 1984) show that the radiation can be highly elliptically polarized and highly variable, if not stochastic. In many cases, the mean of the polarization position angle varies in an S-shaped pattern across the pulse. But histograms of position angle created from the single pulse observations show the angles follow the pattern in two parallel paths separated by about 90 degrees (Stinebring et al. 1984). Furthermore, histograms of fractional linear polarization show that the radiation is significantly depolarized at pulse locations where these orthogonally polarized (OPMs) modes occur. The OPMs are thought to be the natural modes of wave propagation in pulsar magnetospheres (Allen \& Melrose 1982; Barnard \& Arons 1986). The narrow bandwidths and short sampling intervals used in single pulse observations cause the instrumental noise in these observations to be large. The narrow bandwidths are used to overcome pulse smearing effects caused by the dispersion measure of, and multipath scattering in, the interstellar medium. The short sampling intervals, typically of order 100us, are needed to adequately resolve the short duration radio pulse. The combination of the stochastic nature of the intrinsic emission and the high instrumental noise suggests that a statistical approach is needed to analyze the single pulse data. Most results from single pulse polarization observations have been reported as histograms of fractional linear polarization, fractional circular polarization, and polarization position angle (Backer \& Rankin 1980; Stinebring et al. 1984). While these display methods are extremely useful, they do not provide a complete picture of pulsar polarization because they force a separate interpretation of the circular and linear polarization, instead of a combined one as the observed elliptical polarization of the radiation would suggest. A complete, three-dimensional view of the polarization can be made by plotting the polarization measurements from a specific pulse location in the Poincar\'e sphere and projecting the result in two dimensions. The projections show how the orientation of the polarization vector fluctuates on the Poincar\'e sphere and reveal a wide variety of quasi-organized patterns. For example, in the cone emission at the edges of the pulse in PSR B0329+54 (Edwards \& Stappers 2004), the patterns consist of two clusters of data points, each in a separate hemisphere of the Poincar\'e sphere. In the precursor to the pulsar's central core component, the pattern is a single cluster of data points. Within the pulsar's core emission at the center of the pulse, one of the two clusters seen in the cone emission stretches into an ellipse or bar, while the other spreads into an intriguing partial annulus. The signatures of these patterns are not apparent in histograms of fractional polarization or position angle, emphazing the benefit of analyzing the Stokes parameters together. Any viable model of pulsar polarization must be able to replicate the observed patterns in addition to the histograms of fractional polarization.

A statistical model has been developed for the polarization of pulsar radio emission. The model can explain a wide variety of polarization patterns observed in the radio emission. The observations are thus consistent with the model's hypothesis that the polarization of the radiation is determined by the simultaneous interaction of two, highly polarized, orthogonal modes. The analysis of the polarization data shows that polarization signatures of physical processes can become apparent when the Stokes parameters are analyzed together, instead of separately. An interpretation of the model's assumptions and its application to the observations suggest that generalized Faraday rotation may be operative in pulsar magnetospheres. The model shows, in a rigorous way, that polarization measurements follow the statistics of direction.

1012.5434

1012

1012.0816_arXiv.txt

{ We introduce a new, very deep neutral hydrogen (\HI) survey being performed with the Westerbork Synthesis Radio Telescope (WSRT). The Westerbork Hydrogen Accretion in LOcal GAlaxieS (HALOGAS) Survey is producing an archive of some of the most sensitive \HI\ observations available, on the angular scales which are most useful for studying faint, diffuse gas in and around nearby galaxies. The survey data are being used to perform careful modeling of the galaxies, characterizing their gas content, morphology, and kinematics, with the primary goal of revealing the global characteristics of cold gas accretion onto spiral galaxies in the local Universe. In this paper, we describe the survey sample selection, the data acquisition, reduction, and analysis, and present the data products obtained during our pilot program, which consists of UGC~2082, NGC~672, NGC~925, and NGC~4565. The observations reveal a first glimpse of the picture that the full HALOGAS project aims to illuminate: the properties of accreting \HI\ in different types of spirals, and across a range of galactic environments. None of the pilot survey galaxies hosts an \HI\ halo of the scale of NGC~891, but all show varying indications of halo gas features. We compare the properties of detected features in the pilot survey galaxies with their global characteristics, and discuss similarities and differences with NGC~891 and NGC~2403.}

\label{section:discussion} The HALOGAS pilot survey, while still limited in the number of targets, nevertheless covers a substantial range of galactic properties. From the perspective of environmental properties, we have an apparently isolated galaxy (UGC~2082), two targets that have small, gas-rich companions (NGC~925 and NGC~4565), and one galaxy in an interacting pair (NGC~672). We also sample about two orders of magnitude in star formation rate, and nearly a factor of three in rotation speed, with this pilot survey. While it is clear that none of the pilot survey galaxies have an \HI\ halo of the scale seen in NGC~891, we find that all of the pilot survey targets show varying indications of cold gas in their halos. The lack of a prominent halo in UGC~2082 may not be a surprising result: the star formation rate is low, and the galaxy is isolated. Apart from the possibility of significant primordial accretion, a ready source for creating a gaseous halo does not seem to be present in this target. Nonetheless, we find several distinct extraplanar features (these are noted in \S\,\ref{subsection:u2082} and are visible in Figures \ref{figure:u2082} and \ref{figure:pvdiagram_u2082}). These clouds may be a sign of cold gas accretion. Two of the pilot galaxies that are located in denser environments have a much less regular \HI\ structure. Clearly, the case of NGC~672 is a complicated one -- the strong interaction with its companion IC~1727 has spread \HI\ emission throughout the system. This is a dramatic case of interaction-induced accretion. Another interacting system but at a later stage, NGC~925 seems to have interacted with a gas-rich companion sometime in its recent past, resulting in stripped material surrounding the main disk. Whether the culprit is the faint blob at the southern edge of the galaxy, the object found to the north (as described in \S\,\ref{subsection:n0925}), or some other entity, is not yet clear. NGC~925 also seems to possess \HI\ at lagging velocities with respect to the main disk. The anomalous component in NGC~925 has a similar mass and kinematic lag when compared to the same feature in NGC~2403. Another target to compare with existing observations is the large edge-on NGC~4565, which does not show evidence of a prominent halo component. Can this lack be explained? In comparison with NGC~891, the differences in the global parameters of the two galaxies all tend to be in the sense in which one would indeed expect a larger halo in NGC~891, if gaseous halos are induced primarily by star formation-driven galactic fountain flows. Under that assumption, the most important properties of a galaxy for determining whether a halo is formed should be (1) the SF energy injection rate \citep[per unit disk area; see e.g.][]{dahlem_etal_1995,rossa_dettmar_2003a}, and (2) the gravitational potential (traced by $v_{\mathrm{rot}}^2$). With NGC~891's larger SFR, smaller disk area (at the distances listed in Table \ref{table:sample}), and lower rotational velocity, it would seem that it should indeed be more efficient in maintaining a gaseous halo. Per unit disk area, more star formation energy is available, and gravity is less effective at preventing material from moving away from the disk midplane. That NGC~925 and NGC~2403 have similar anomalous gas features, as well as similar SFR and mass, would also be consistent with such a picture of gaseous halo formation. Do the other HALOGAS targets follow the same trend with the global galaxy parameters? Though it is intriguing, the number of galaxies which have been observed with enough sensitivity to distinguish a population of halo gas remains rather low, so that general statements are still uncertain. This initial result from our pilot survey underlines the need for a suitably large survey, in order to trace a range of variables in the population of nearby galaxies. The full HALOGAS survey will provide such a large sample.

1012.0816

1012

1012.5328_arXiv.txt

Recent radio observations reveal the existence of mini radio lobes in active galaxies with their scales of $\sim 10~{\rm pc}$. The lobes are expected to be filled with shock accelerated electrons and protons. In this work, we examine the photon spectra from the mini lobes, properly taking the hadronic processes into account. We find that the resultant broadband spectra contain the two distinct hadronic bumps in $\gamma$-ray bands, i.e., the proton synchrotron bump at $\sim$ MeV and the synchrotron bump at $\sim$ GeV due to the secondary electrons/positrons produced via photo-pion cascade. Especially when the duration of particle injection is shorter than the lobe age, radio-dark $\gamma$-ray lobes are predicted. The existence of the $\gamma$-ray lobes could be testable with the future TeV-$\gamma$ telescope {\it CTA}.

Thanks to the progress of VLBI (Very Long Baseline Interferometry) observations, compact radio lobes with a linear size $LS\sim 1~{\rm kpc}$, defined as a projected length from the core to the lobe, have been discovered (e.g., Fanti et al. 1995; Readhead et al. 1996; O'Dea \& Baum 1997). Further VLBI observations recently reveal the existence of very small radio lobes with $LS\sim 10~{\rm pc}$ among two samples of compact radio sources. The sample of compact radio sources with their spectral peaks higher than $\sim5~{\rm GHz}$ is termed as high frequency peakers (HFPs). Some of them turn out to be mini lobes based on their morphologies and non-variabilities (e.g., Orienti et al. 2007; Orienti \& Dallacasa 2008). The sample of compact radio sources at low-redshift ($z<0.16$) is selected and called as CORALZ. The VLBI observations of CORALZ show that some of them are also found as mini radio lobes (Snellen et al. 2004; de Vries et al. 2009). Apart from the above two cases, recurrent radio sources are also known to possess mini lobes inside large lobes (e.g., 3C 84, Walker et al. 2000). These mini radio lobes are young and their ages are typically estimated as $t_{\rm age}\sim 10^{2-3}~{\rm yrs}$ (e.g., Fanti 2009, Giroletti \& Polatidis 2009). High-energy emission of the mini radio lobes have been theoretically explored by some authors (Stawarz et al. 2008; Kino et al. 2007, 2009). However, previous work has focused on leptonic processes and little is known about hadronic processes in the lobes. In this Letter, we indicate that the hot spots in the mini radio lobe can be a plausible site of proton acceleration. Therefore, high energy protons are naturally expected in the mini lobes. In order to constrain the amount of high energy protons, next we examine the predicted photon spectra. Because of their smallness, observed mini radio lobes generally have dense synchrotron photon fields. Furthermore, the mini lobes are close to the AGN core and illuminated by the core emission. Therefore $p\gamma$ interaction is inevitable for the mini lobes. We calculate the photon spectra from the mini lobes taking into account the hadronic processes and show that the spectra are useful for constraining the amount of protons via $\gamma$-ray emission features. The predicted spectra will be useful for testing whether cosmic ray acceleration indeed takes place in mini lobes.

In this work, we point out that the hot spots in mini lobes are feasible sites of proton acceleration. Next, the expected photon spectra of the mini lobes including the hadronic processes are explored. Summary and discussions are as follows. \begin{enumerate} \item For bright lobes with $L_{e}\sim L_{\rm syn}\sim 10^{45}~{\rm erg~s^{-1}}$, the predicted high energy emission is detectable with the current $\gamma$-ray telescopes. However, it is overwhelmed by the leptonic inverse Compton component, and it seems hard to test whether the emission is of hadronic- or leptonic-origin. For $L_{e}\sim L_{\rm syn}\le 10^{43}~{\rm erg~s^{-1}}$, proton synchrotron bump appears at $\sim$ MeV and the synchrotron emission from the secondary electrons/positrons generated by the $\mu$- and $\pi$- decays emerges at $\sim$ GeV. The two distinctive bumps in $\gamma$-ray domain are of hadronic origin. \item The case of the short term particle injection is examined. It may be realized when jets have intermittent activities. Typically, primary accelerated electrons have been cooled down but protons have not. Then, the predicted emission is purely hadronic and the hadronic bumps are clearly seen. Importantly, the high energy tail of the GeV bump is detectable by {\it The Cherenkov Telescope Array (CTA)} (http://www.cta-observatory.org/). These sources may be identified as radio dark $\gamma$-ray lobes. \item The predicted X-ray flux is well above the detection limit of $XMM$. The observations actually show bright $X$-ray emission. This is traditionally interpreted as thermal radiation from the accretion disk and the possibility of lobe emission has been alternatively indicated (Ostorero et al. 2010 and reference therein). In any case, X-ray emission is likely composed of various different components. Therefore, it seems difficult to extract the hadronic component from the X-ray band. \item We comment on the importance of larger $n_{\rm ext}$ it may lead to frre-free absorption (FFA) in radio band (e.g., Begelman 1999; Bicknell 2003; Stawarz et al. 2008). Actually, the low frequency turnover in the radio spectra of some sources are indeed reproduced by FFA and not by SSA (e.g., OQ 208, Kameno et al. 2000; 0108+388, Marr et al. 2001). It is clear that dense emvironments lead to an effective proton-proton collision. We will examine it in the future. \item We add a comment on the recent VLBI observation of mini lobe 3C 84. It shows the outburst around 2005 and a new component smaller than $1~{\rm pc}$ emerges (Nagai et al. 2010). Fermi/LAT also detect GeV $\gamma$-ray emission from it (Abdo et al. 2009). A future collaboration with Space VLBI project VSOP-2 with high angular resolution (http://www.vsop.isas.jaxa.jp/vsop2/) and {\it CTA} would provide us valuable constraints on the hadronic model. Theoretically, we plan to conduct studies with smaller $R_{\rm lobe}$ in our future work.

1012.5328

1012

1012.5602_arXiv.txt

The Sculptor and Carina Dwarf spheroidal galaxies were observed with the H.E.S.S. Cherenkov telescope array between January 2008 and December 2009. The data sets consist of a total of 11.8 and 14.8 hours of high quality data, respectively. No gamma-ray signal was detected at the nominal positions of these galaxies above 220 GeV and 320 GeV, respectively. Upper limits on the gamma-ray fluxes at 95\%~C.L. assuming two forms for the spectral energy distribution (a power law shape and one derived from dark matter annihilation) are obtained at the level of 10$^{-13}$ to 10$^{-12}$~cm$^{-2}$s$^{-1}$ in the TeV range. Constraints on the velocity weighted dark matter particle annihilation cross section for both Sculptor and Carina dwarf galaxies range from $\left\langle \sigma v \right\rangle\sim 10^{-21}$~cm$^3$s$^{-1}$ down to $\left\langle \sigma v \right\rangle\sim 10^{-22}$~cm$^3$s$^{-1}$ depending on the dark matter halo model used. Possible enhancements of the gamma-ray flux are studied: the Sommerfeld effect, which is found to exclude some dark matter particle masses, the internal Bremsstrahlung and clumps in the dark-matter halo distributions.

1012.5602

1012

1012.0205_arXiv.txt

{ High-velocity stars in the Galactic halo, e.g. the so-called hyper-velocity stars (HVS), are important tracers of the properties of the dark matter halo, in particular its mass. }{ A search for the fastest stars among hot subdwarfs (sdB) in the halo is carried out to identify HVS, unbound to the Galaxy, and bound population II stars in order to derive a lower limit to the halo mass. }{ Based on the SDSS DR6 spectral database we selected stars with high rest-frame velocities. These radial velocity measurements were verified at several telescopes to exclude radial velocity variable stars. Out of 88 stars observed in the follow-up campaign 39 stars were found to have constant radial velocities. For twelve of them we measured a proper motion significantly different from zero and obtained spectroscopic distances from quantitative spectral analysis to construct the full 6D phase space information for a kinematical study. }{ All but one programme sdBs show halo characteristics, but can be distinguished into two kinematical groups, one (G1) with low Galactic rotation typical of halo stars and a second one (G2) with rapid retrograde motion. We also investigate the possibility that the programme stars are not genuine halo stars but ejected from the Galactic disc or bulge. The G1 objects crossed the Galactic plane in the central bulge, whereas the G2 stars did in the outer Galactic disc. J1211+1437 (G2) is a HVS candidate, as it is unbound to the Galaxy if the standard Galactic potential is adopted. }{ We conclude that in the ejection scenario G1 stars might have been formed via the slingshot mechanism that invokes acceleration by tidal interaction of a binary with the central supermassive black hole. The G2 stars, however, would originate in the outskirts of the Galactic disc and not in the central bulge. J1211+1437 is the first unbound subdwarf B star, for which we can rule out the slingshot mechanism. Alternatively, we may assume that the stars are old population II stars and therefore have to be bound. Then the kinematics of J1211+1437 set a lower limit of $2\times 10^{12}$ M$_\odot$ to the mass of the Galactic dark matter halo. }

\label{sec:intro} The properties of the dark matter halo are important to understand how the Galaxy formed and evolved. Observations of halo stars put constraints on theoretical models of halo formation \citep[e.g.][]{1996ApJ...462..563N}. Large surveys, such as the Sloan Digital Sky Survey \citep[SDSS,][]{2000AJ....120.1579Y} and the RAdial Velocity Experiment \citep[\emph{RAVE},][]{2006AJ....132.1645S}, provide large numbers of stars to trace the halo properties, such as the total mass of the halo. Globular clusters, satellite galaxies, as well as large samples of halo stars, respectively, have been used to estimate the halo mass. Actually only the objects with the most extreme velocities provide tight constraints and, hence, the mass estimates depend mostly on them \citep{2003A&A...397..899S,2007MNRAS.379..755S}. A halo mass of about $2\times10^{12}$\,\msun\ was favoured in earlier investigations \citep{1999MNRAS.310..645W,2003A&A...397..899S}, while more recent studies prefer lower masses of about half that value \citep{2005MNRAS.364..433B,2007MNRAS.379..755S,2008ApJ...684.1143X}. The hyper-velocity stars \citep[HVS, ][]{2005ApJ...622L..33B,2005A&A...444L..61H,2005ApJ...634L.181E} are the fastest moving stars known in the halo. Their supposed place of origin is the Galactic centre, where they have been suggested to be accelerated by tidal interactions of a binary star with the super-massive black hole \citep[SMBH, ][]{1988Natur.331..687H}. Whether a HVS can in fact escape from the Galaxy or not depends on the halo mass \citep{2009ApJ...691L..63A}. Kinematical studies of the hyper-velocity stars were based on their radial velocities only. Recently, \cite{2009A&A...507L..37T} were able to measure proper motions of an A-type HVS and study its 3-D kinematics to trace its place of birth in the Galactic disc. They found it to originate far from the Galactic centre, thereby challenging the SMBH-slingshot mechanism of \cite{1988Natur.331..687H}. Hence \citet{2009A&A...507L..37T} suggested a runaway mechanism for the star's formation. Further evidence that such a mechanism works comes from two similar studies of the hyper run-away stars HD 217791 \citep{2008A&A...483L..21H} and HIP~60350 \citep{2010ApJ...711..138I}, which were also found to originate in the outer rim of the Galactic disc nowhere near the Galactic centre. While most of the 17 HVS known today \citep{2009ApJ...690.1639B, 2009A&A...507L..37T} are early-type main-sequence stars, there is just one evolved low-mass star, US~708, a hot subdwarf star of spectral type sdO \citep{2005A&A...444L..61H}. Most of the previous studies of halo stars to constrain the dark matter properties are hampered by the lack of proper motion measurements. Hence they had to rely substantially on radial velocity distributions. In such cases only four coordinates (i.e. two position values, distance and radial velocity, RV) of the 6D phase space are determined and the missing proper motion components are handled in a statistical approach. In the presently most extensive study \citet{2008ApJ...684.1143X} measured radial velocities for more than 10,000 blue halo stars from the SDSS and classified their sample as a mix of blue horizontal branch (BHB) stars, blue stragglers and main-sequence stars with effective temperatures roughly between 7,000 and 10,000\,K according to their colours. \citet{2008ApJ...684.1143X} selected 2400 blue horizontal-branch stars to estimate the halo mass out to 60~kpc to be $1.0\times10^{12}$ M$_\odot$ using a halo model of \cite{1997ApJ...490..493N}. For one star from that sample \cite{2010ApJ...718...37P} were able to obtain proper motion and carry out a detailed kinematic analysis, which revealed an inbound Population~II horizontal branch star with a Galactic rest-frame (GRF) velocity of $\sim$700\,\kms\ at its current position. This makes it the fastest halo star known, and provided a lower limit of $1.7\times10^{12}$ M$_\odot$ for the total halo mass of the Galaxy, significantly exceeding the value determined by \citet{2008ApJ...684.1143X}. This example shows that it is rewarding to study the kinematics of additional stars in the halo and to consider classes of stars other than BHB stars, as well. Of course, the Galactic halo hosts a plethora of white dwarfs \citep{2006ApJS..167...40E}. However, they are so faint that they can be analysed in the solar neighbourhood only. Another group of evolved low mass stars are the hot subdwarf stars (sdB, sdO) that dominate the population of faint blue stars at high Galactic latitudes to visual magnitudes of about V=18 \citep{1986ApJS...61..305G}. They are considered to be helium core burning stars with very thin ($<$0.02\,$M_\odot$) inert hydrogen envelopes and masses around 0.5~\msun. Following ideas outlined by \cite{1986A&A...155...33H}, the sdBs can be identified with models for extreme horizontal branch (EHB) stars. An EHB star bears great resemblance to a helium main-sequence star of half a solar mass and it should evolve similarly, i.e. directly to the white dwarf cooling sequence, bypassing a second giant phase \citep[for a review see][]{2009ARA&A..47..211H}. For the formation of subdwarf B stars three scenarios are discussed by \cite{2003MNRAS.341..669H}: common envelope ejection, stable Roche lobe overflow (RLOF), and the merger of two helium white dwarfs. Some alternate scenarios for the formation of single sdB stars are reviewed by \cite{2009CoAst.159...75O}. Hot subdwarf stars exist in the field of the Galaxy but also in globular clusters, in the Galactic bulge and have even been resolved in the elliptical galaxy M~32 \citep{2008ApJ...682..319B}. Kinematical studies \citep{2004A&A...414..181A, 2008ASPC..391..257N} indicate that they occur in all stellar populations of the Galaxy. However, very little is known about the halo population of hot subdwarfs except those in globular clusters. Some high-velocity hot subdwarfs have attracted interest because of their high radial velocities, most notably, the sdO star US~708, whose radial velocity in the rest-frame was measured at 751 \kms \citep{2005A&A...444L..61H} -- the second HVS star discovered. Unfortunately we cannot deduce the origin of the star, as we lack a reliable proper motion measurement. Motivated by the discovery of US~708, we embarked on a project to identify a sample of population II hot subdwarfs and study their kinematics from radial velocity and proper motion. We make use of the MUCHFUSS survey \citep{geier_TARGET}, which searches for close binaries with high radial-velocity variations. The search strategy also provides targets that are not close binaries but travel through space at high RV without variations. These stars are the targets of our investigation. Accordingly we entitled our project Hyper-MUCHFUSS as we provide an extension of MUCHFUSS. The paper is organised as follows. In Sect.~\ref{sec:survey} we introduce our survey for HVS and in Sect.~\ref{sec:PM} our sophisticated proper motion measurement method. The kinematical analysis techniques are shown in Sect.~\ref{sec:DIST_KINE}. In Sect.~\ref{sec:results} we present our results and summarise and conclude in Sect.~\ref{sec:conclu}.

1012.0205

1012

1012.2807_arXiv.txt

We present near infrared (NIR) IRTF/SpeX spectra of the intermediate-age galaxy M32 and the post-starburst galaxy NGC 5102. We show that features from thermally-pulsing asymptotic giant branch (TP-AGB) and main sequence turn-off (MSTO) stars yield similar ages to those derived from optical spectra. The TP-AGB can dominate the NIR flux of a coeval stellar population between $\sim$0.1 and $\sim$2~Gyr, and the strong features of (especially C-rich) TP-AGB stars are useful chronometers in integrated light studies. Likewise, the Paschen series in MSTO stars is stongly dependent on age and is an indicator of a young stellar component in integrated spectra. We define four NIR spectroscopic indices to measure the strength of absorption features from both C-rich TP-AGB stars and hydrogen features in main sequence stars, in a preliminary effort to construct a robust chronometer that probes the contributions from stars in different evolutionary phases. By comparing the values of the indices measured in M32 and NGC~5102 to those in the \citet{m05} stellar population synthesis models for various ages and metallicities, we show that model predictions for the ages of the nuclei of M32 and NGC~5102 agree with previous results obtained from integrated optical spectroscopy and CMD analysis of the giant branches. The indices discriminate between an intermediate age population of $\sim$3--4~Gyr, a younger population of $\lesssim$1~Gyr, and can also detect the signatures of very young $\lesssim$100~Myr populations.

\label{sec:int} Integrated spectroscopy of galaxies has been performed primarily in the optical because much of the (blue) optical light comes from the well-understood main sequence turnoff stars. As one pushes into the near infrared (NIR), the integrated light of a stellar population is dominated by very luminous stars in later stages of stellar evolution, such as thermally pulsing asymptotic giant branch (TP-AGB) stars. Beginning with \citet{renzini81} and continuing with, e.g.,\ \citet{frogel90}, \citet{bressan98}, \citet{m98}, \citet{lancon99}, \citet{lancon02}, it has become clear that the unique features and high luminosity of TP-AGB stars profoundly affect the NIR integrated light of stellar populations. Specifically, for stellar populations with ages $\sim$100~Myr to $\sim$1~Gyr, the TP-AGB can contribute up to $\sim$60\% of the K-band flux \citep[depending on metallicity, see][]{m05}. Strong molecular features, from CN, CO, C$_2$, and H$_2$O, in TP-AGB stellar atmospheres should appear in the integrated NIR spectrum of populations in this age range. Unfortunately, relatively poorly understood processes such as mass loss, convective transport of heavy elements to the surface, and thermal pulsation influence the late stages of stellar evolution and thus predictions for the integrated spectrum. Furthermore, the ratio of C-rich to O-rich AGB stars in a population is known to depend strongly on metallicity, meaning that the contribution from TP-AGB stars can vary drastically not only with age but also composition \citep{mouhcine03}. However, observational work has constrained these effects on the integrated light of stellar populations with observations of Magellanic Cloud and Galactic clusters, and the predictive potential of NIR spectroscopic features has been demonstrated \citep[e.g.,\ ][]{lancon99,mouhcine02,m05}. In short, TP-AGB features should be traceable in galaxies whose NIR light is dominated by stars in the $\sim$100~Myr to $\sim$1~Gyr age range. For example, \citet{mouhcine02a} have detected the presence of TP-AGB stars in the NIR spectrum of a young star cluster in the galaxy NGC 7252, and have verified that the predicted features expected in a young population ($\sim$300~Myr) are present in the NIR spectrum. This successful detection of TP-AGB features in NGC~7252 thus motivates additional NIR spectroscopy of bright, nearby galaxies with young populations to further test the use of NIR spectral indicators as chronometers. We also note that while light from the TP-AGB can dominate the NIR flux, the main sequence turn-off (MSTO) also contributes an appreciable amount to the integrated light, and hydrogen features (specifically the Paschen series) due to MSTO stars will be visible. To investigate the behavior of the TP-AGB and MSTO features as a function of both age and metallicity we consulted the \citet[][hereafter M05]{m05} stellar population synthesis models, which include a careful treatment of the TP-AGB. In this Letter we use NIR spectra of two galaxies with very different star formation histories to test if population synthesis models of TP-AGB and MSTO features can differentiate between young and intermediate-age populations. Because its high nuclear surface brightness and proximity allow for both high SNR spectroscopy and resolved stellar photometry of the giant branch, the intermediate age compact elliptical galaxy M32 has been extensively studied. Recent work has included integrated optical spectroscopy of the nucleus and extranuclear region out to the half-light radius $R_e$ \citep{delburgo01,worthey04,rose05,coelho09}, and NIR imaging of the giant branches \citep{davidge07}. These studies find that the nuclear light is dominated by an approximately solar metallicity intermediate-age population ($\sim$3 -- 4~Gyr). Because its stellar content has been constrained by several established age-dating techniques, M32 is ideal for testing other predictive techniques. Likewise, the blue S0 galaxy NGC~5102 is close and bright enough for high SNR integrated spectroscopy plus resolved CMD analysis of its brightest stars. It is classified as a post-starburst (PSB) galaxy because its strong Balmer-line absorption and absence of emission indicates that star formation terminated recently. Studies by, e.g., \citet{deharveng97,davidge08} have indicated that the nuclear region of NGC 5102 has undergone a termination of star formation within the past $\sim$10--100~Myr, after a several hundred Myr period of activity, guaranteeing that there is an appreciable young population with a mean age of a few 100 Myr. The intermediate age stars in M32 and the younger stars in NGC~5102 are a useful pairing because most of the NIR light emitted by these two galaxies comes from stars in different evolutionary stages: the younger population will be dominated by the TP-AGB and a blue MSTO, while in the intermediate age population most of the light comes from the red giant branch and an older MSTO (see Fig.~13 in M05). In this Letter we define four age-sensitive NIR spectroscopic indices, two of which measure previously defined TP-AGB features, and two that measure Paschen series absorption lines, and compare the observed values in M32 and NGC~5102 to predictions of the M05 stellar population models. We demonstrate that the derived ages from the NIR indices agree closely with previously published ages determined from optical spectroscopy, and thus provide evidence to support the reliability of NIR spectral signatures as a useful chronometer for galaxy evolution.

We present an analysis of integrated NIR SpeX SXD spectra of the nuclear regions of M32 and NGC~5102, and show that mean ages determined from spectroscopic indices agree with previous studies, with the important result that this method can differentiate between young and intermediate-age populations. The indices probe contributions from two different stellar evolutionary phases, the TP-AGB and the MSTO, which effectively provides two independent chronometers. Galaxy ages are derived by comparing index values to those measured in M05 model SSPs, indicating an accurate treatment of the TP-AGB in the models. This method of defining NIR indices is the first step towards a robust NIR spectroscopic age-dating technique, which will be particularly useful when applied to high redshift spectroscopic surveys undertaken by the next generation of infrared observatories.

1012.2807

1012

1012.0519_arXiv.txt

Tabulated rates for astrophysical photodisintegration reactions make use of Boltzmann statistics for the photons involved as well as the interacting nuclei. Here we derive analytic corrections for the Planck-spectrum quantum statistics of the photon energy distribution. These corrections can be deduced directly from the detailed-balance condition without the assumption of equilibrium as long as the photons are represented by a Planck spectrum. Moreover we show that these corrections affect not only the photodisintegration rates but also modify the conditions of nuclear statistical equilibrium as represented in the Saha equation. We deduce new analytic corrections to the classical Maxwell-Boltzmann statistics which can easily be added to the reverse reaction rates of existing reaction network tabulations. We show that the effects of quantum statistics, though generally quite small, always tend to speed up photodisintegration rates and are largest for nuclei and environments for which $Q/kT \sim 1$. As an illustration, we examine possible effects of these corrections on the $r$-process, the $rp$-process, explosive silicon burning, the $\gamma$-process and big bang nucleosynthesis. We find that in most cases one is quite justified in neglecting these corrections. The correction is largest for reactions near the drip line for an $r$-process with very high neutron density, or an $rp$-process at high-temperature.

The capture of a projectile particle by a nucleus followed by the emission of a photon is called radiative capture. The inverse process is called photodisintegration. Both types of reactions play important roles in stellar and big bang nucleosynthesis \citep*{rol90,smi93,wag67,fow67,cla68,Iliadis07}. Photodisintegration in astrophysical environments often involves a thermal distribution of photons that excite nuclei above the particle emission threshold. Above a temperature of $T \sim 10^9$ K, photodissociation can become a dominant process in the reaction flow and the photo-ejected nucleons can be captured by other nuclei leading to photodisintegration rearrangement as can happen for example during core or explosive oxygen or silicon burning \citep{cla68, Iliadis07}. In determining the rate for photodisintegration reactions, however, one should take into account various factors arising from dealing with photons in a two body problem. One of them is the fact that photons are massless bosons, and hence, obey Planckian statistics. Historically, photodisintegration reactions have been treated with Maxwell-Boltzmann statistics, both because it is usually an excellent approximation (cf. \cite{Rauscher95,Iliadis07}) and because this assumption simplifies the determination of the photodisintegration rates. However, low-energy photons are more correctly represented by a Planck distribution. As discussed below, if the ratio of the capture $Q$-value to the temperature $Q/kT$ is large then the use of a Maxwell-Boltzmann distribution for the photons is a good approximation. On the other hand, for $Q/kT \sim 1$ (e.g., if one is interested in the photo-ejection of loosely bound particles from nuclei), then the effects of quantum statistics become more relevant. We show here that modified thermal photodisintegration rates can be written in the form of a small analytic correction to the tabulated reaction rates obtained with Maxwell-Boltzmann statistics. Furthermore, these correction factors are nearly independent of the nuclear cross sections and to leading order only depend upon the reaction $Q$-value, the Gamow energy (for the charged-particle nonresonant part), and the resonance energy (for the resonant part). \subsection{Photonuclear Reactions} For a reaction (not necessarily in equilibrium) of the form \begin{equation}\label{Reaction} 1+2 \rightarrow 3+\gamma~~, ~~3+\gamma \rightarrow 1 + 2~~, \end{equation} the forward and reverse reaction rates are given by \begin{eqnarray} r_{12}&=&n_{1}n_{2}\langle\sigma v\rangle_{12} = n_{1} \lambda_{12} \\ r_{\gamma3}& = &n_{3}n_{\gamma}\langle\sigma c\rangle_{\gamma3}= n_{3} \lambda_{\gamma 3}~~. \label{eq:Rates} \end{eqnarray} Here, $\langle\sigma v\rangle_{12}$ denotes the thermally averaged reaction rate per particle pair for the capture reaction and is given by an integration over an appropriate velocity distribution $\phi(v)$, \begin{equation} \langle\sigma v\rangle_{12}=\int\sigma_{12}\left(v\right)\: v\:\phi(v)\: d^{3}v~~, \end{equation} where $v$ denotes the relative velocity of the nuclei $1$ and $2$. For most cases of interest in astrophysics, the massive interacting particles are non-degenerate (i.e., dilute) and non-relativistic (i.e., their rest mass energy is large compared to $kT$). Hence, one can use a Maxwell-Boltzmann velocity distribution with Newtonian kinetic energy because \begin{equation} \frac{1}{e^{\sqrt{p^{2}c^{2}+m^{2}c^{4}}/kT}\pm1}\simeq e^{-\sqrt{p^{2}c^{2}+m^{2}c^{4}}/kT}\propto e^{-\frac{mv^{2}}{2kT}}~~, \end{equation} where the exponential factor involving the rest mass energy drops out with a proper normalization of the distribution function. Also note that when both particles $1$ and $2$ obey Maxwell-Boltzmann statistics, so does their relative velocity \citep{cla68}. Hence, the thermally averaged capture rate per particle pair is given by \begin{equation} \langle\sigma v\rangle_{12}=\sqrt{\frac{8}{\pi\mu}}\left(kT\right)^{-\frac{3}{2}}\int_{0}^{\infty}\sigma_{12}\left(E\right)\: e^{-E/kT}\: E\: dE~~,\label{eq:Forward Reaction Rate} \end{equation} where $\mu$ denotes the reduced mass for the particles $1$ and $2$. In the photodisintegration reaction rate, however, the relative velocity of the photon with respect to the target nucleus is always the speed of light $c$. This eliminates any dependence of the reaction rate on the velocity distribution of the target nuclei \citep{thi98}. Therefore, the photonuclear reaction rate $\langle\sigma c\rangle_{\gamma3}$ becomes an integral of the reaction cross section over a Planck energy distribution for the photons: \begin{equation} \langle\sigma c\rangle_{\gamma3}=\frac{1}{2 \zeta(3) \left(kT\right)^{3}} \int_{\mbox{\footnotesize max}(0,Q)}^{\infty}\sigma_{\gamma3}\left(E_{\gamma}\right)\: c\:\frac{1}{e^{E_{\gamma}/kT}-1}\: E_{\gamma}^{2}\: dE_{\gamma}~~. \label{eq:Inverse Reaction Rate} \end{equation} Using the fact that for a Planck distribution \begin{equation} n_\gamma = {16 \pi} { \zeta(3)}\biggl(\frac {kT}{h c}\biggr)^{3}~~, \end{equation} we can equivalently write Eq.~(\ref{eq:Inverse Reaction Rate}) in more familiar form \begin{equation} \lambda_{\gamma 3}=\frac{8 \pi}{(h c)^{3}} \int_{\mbox{\footnotesize max}(0,Q)}^{\infty}\sigma_{\gamma3}\left(E_{\gamma}\right)\: c\:\frac{1}{e^{E_{\gamma}/kT}-1}\: E_{\gamma}^{2}\: dE_{\gamma}~~, \label{eq:Inverse lambda} \end{equation} where $\zeta(3) = 1.20206$ is the Riemann zeta function and $E_{\gamma}$ denotes the photon energy. The integration threshold is the $Q$-value of the capture reaction (see Figure \ref{fig:Energy Figure}) or zero in the case of negative $Q$. \placefigure{fig:Energy Figure} \subsection{Detailed Balance Condition} It is difficult to determine the cross section $\sigma_{\gamma3}$ directly from experiment. However, the interaction between photons and matter is very weak $(e^{2}/\hbar c\ll1)$ so that the reaction can be treated with first order perturbation theory. In this case, the transition probabilities become proportional to the matrix elements of the perturbing Hamiltonian and the hermiticity of the perturbing Hamiltonian gives rise to a simple relation between the capture and disintegration cross sections. This is known as the \emph{detailed balance equation} \citep{bla91}. For a reaction involving a ground-state to ground-state transition for two nuclei with energy $E$, leading to a gamma ray with energy $E_{\gamma}=E+Q$ this is given by \begin{equation} \sigma_{\gamma3}(E_{\gamma})=\frac{g_{1}g_{2}}{ g_{3}(1 + \delta_{12})}\frac{\mu c^2 E}{E_{\gamma}^{2}}\sigma_{12}(E)~~, \label{eq:Detailed Balance} \end{equation} where $g_{i} = 2 j_i + 1 $ are the spin degeneracy factors for the ground state of the nuclei and the Kronecker delta function accounts for the special case of indistinguishable interacting nuclei. Using this detailed balance equation, the photodisintegration rate for a single-state transition can be related to the forward capture rate. Substituting Eq.~(\ref{eq:Detailed Balance}) into Eq.~(\ref{eq:Inverse Reaction Rate}) and also changing the variable from $E_{\gamma}$ to $E=E_{\gamma}-Q$ in the integration, one can average over the velocity distribution of the ground-state interacting nuclei $\langle\sigma c\rangle_{\gamma3}$ as follows: \begin{equation} \langle\sigma c\rangle_{\gamma3}=\frac{\mu c^{3}}{2 \zeta(3) (kT)^{3}}\frac{g_{1}g_{2}}{g_{3} (1 + \delta_{12})}\int_{0}^{\infty}\sigma_{12}\left(E\right)\frac{1}{e^{\left(E+Q\right)/kT}-1}EdE. \label{eq:Inverse Reaction Rate 2} \end{equation} At this point, one usually introduces the approximation: \begin{equation} e^{\left(E+Q\right)/kT}-1 \approx e^{\left(E+Q\right)/kT}~~. \label{approx} \end{equation} Here, we point out that by inserting this approximation and then correcting for it, Eq.~(\ref{eq:Inverse Reaction Rate 2}) can be rewritten in the following exact form: \begin{equation} \langle\sigma c\rangle_{\gamma3} = \bigl(1+R\bigr) {\langle\sigma v\rangle_{12}} \biggl( \frac{ \sqrt{2 \pi}}{8 \zeta(3)} \biggr) \frac{g_{1}g_{2}}{g_{3} (1 + \delta_{12})} \left(\frac{\mu c^{2}}{kT}\right)^{3/2}e^{-Q/kT}~, \label{eq:Ratio of Reaction Rates 3} \end{equation} where $R$ is a small and dimensionless number which is formally given by \begin{equation} 1 + R=\biggl[\frac{\int_{0}^{\infty}\sigma_{12}\left(E\right)(e^{\left(E+Q\right)/kT}-1)^{-1}EdE}{\int_{0}^{\infty}\sigma_{12}\left(E\right)e^{-\left(E+Q\right)/kT}EdE}\biggr]~~. \label{eq:R-Formal} \end{equation} \subsection{Thermal Population of Excited States} The generalization of Eq. (\ref{eq:Ratio of Reaction Rates 3}) to the average over thermally populated states among the initial and final nuclei is straightforward \citep{cla68, Iliadis07}. One must first replace the ground state (g.s.) to g.s.~forward reaction cross section $\sigma_{1 2}$ with a weighted average over the thermal population of states $\mu$ in the target nucleus 1, and also sum over all final states in product nucleus $3$. (Note that we only consider light particle $(p,\gamma)$, $(n,\gamma)$, or $(\alpha,\gamma)$, reactions for which we can ignore their excitation.) Thus, the effective stellar thermal forward rate becomes \begin{equation} \langle \sigma v \rangle_{1 2}^* = \frac{\sum_\mu g_{1 \mu} e^{-E_{1 \mu}/kT} \sum_\nu \langle \sigma v \rangle_{1 2}^{\mu \rightarrow \nu}}{\sum_\mu g_{1 \mu} e^{-E_{1 \mu}/kT} }~~, \end{equation} which can also be written as \begin{equation} \langle \sigma v \rangle_{1 2}^* = R_{t t} \langle \sigma v \rangle_{1 2} ~~, \end{equation} where the stellar enhancement factor $R_{t t}$ is defined by \begin{equation} R_{t t} = \frac{\sum_\mu g_{1 \mu} e^{-E_{1 \mu}/kT} \frac{ \sum_\nu \langle \sigma_{1 2} v \rangle_{1 2}^{\mu \rightarrow \nu}}{\sum_\nu \langle \sigma_{1 2} v \rangle_{1 2}^{g.s. \rightarrow \nu}}}{\sum_\mu g_{1 \mu} e^{-E_{1 \mu}/kT} }~~. \end{equation} Usually, tabulated thermonuclear reaction rates are given as the ground state rate and the stellar enhancement factor must be determined from a statistical model calculation as in \cite{Holmes76}, \cite{Woosley78}, \cite{Rauscher00} and \cite{Rauscher04}. The thermally averaged photonuclear rate for a distribution of excited states in initial and final heavy nuclei then becomes \begin{equation} \langle \sigma c \rangle_{3 \gamma}^* = \frac{\sum_\nu g_{3 \nu} e^{-E_{3 \nu}/kT} \sum_\mu \langle \sigma c \rangle_{3 \gamma}^{\nu \rightarrow \mu}}{\sum_\nu g_{3 \nu} e^{-E_{3 \nu}/kT} } ~~. \label{sigmac-star} \end{equation} The generalization of the detailed balance condition of Eq.~(\ref{eq:Ratio of Reaction Rates 3}) is \begin{eqnarray} \langle \sigma c \rangle_{\gamma 3}^{\mu \rightarrow \nu} & = & \bigl(1+R_{\mu \nu} \bigr) {\langle\sigma v \rangle_{12}^{\nu \rightarrow \mu}} \biggl( \frac{ \sqrt{2 \pi}}{8 \zeta(3)} \biggr) \frac{g_{1 \mu}g_{2}}{g_{3 \nu } (1 + \delta_{12})} \nonumber \\ && \times \left(\frac{\mu c^{2}}{kT}\right)^{3/2}e^{-Q_{\mu \nu}/kT}~~, \label{eq:Ratio of Reaction Rates munu} \end{eqnarray} where $R_{\mu \nu}$ denotes the use of $Q_{\mu \nu}$ and $\sigma_{1 2}^{\mu \rightarrow \nu}(E)$ in Eq. (\ref{eq:R-Formal}). Inserting Eq.~(\ref{eq:Ratio of Reaction Rates munu}) into Eq.~(\ref{sigmac-star}) and using the fact that $Q_{\mu \nu} = Q - E_{3 \nu} + E_{1 \mu}$, we can write \begin{eqnarray} \langle \sigma c \rangle_{\gamma 3}^* &=& \bigl(1+R \bigr) {\langle\sigma v\rangle_{12}^*} \nonumber \\ &\times & \biggl( \frac{ \sqrt{2 \pi}}{8 \zeta(3)} \biggr) \left(\frac{\mu c^{2}}{kT}\right)^{3/2} \frac{G_{1}G_{2}}{G_{3} (1 + \delta_{12})}e^{-Q/kT}~~, \label{eq:thermal} \end{eqnarray} where $R$ represents an average correction factor among all thermally populated states. As demonstrated below, $R$ is nearly independent of the detailed nuclear structure. Hence, we can simply utilize the ground-state $Q$-value as a representative average over the distribution of $Q$-values among the thermally populated states. Also, the spin factors above are now replaced by the relevant nuclear partition functions $G_{i}$ : \begin{equation} G_i =\sum_{\alpha}g_{\alpha}e^{-E_{\alpha}/kT}~~, \end{equation} where $\alpha$ denotes the individual states in nucleus $i$. Stellar reaction rate tables are usually listed as functions of temperature $T_9$ in units of $10^9$ K and are given as $\bigl[ N_A \langle \sigma v(T_9) \rangle^* \bigr]$. Thus, we can rewrite Eq.~(\ref{eq:thermal}) as \begin{eqnarray} \label{eq:thermalT9} \lambda_{\gamma 3} &= & \bigl(1+R \bigr) \bigl[N_A \langle\sigma v(T_9) \rangle^* \bigr]_{12} \\ \times 9.8685 \times 10^9 && ({\hat \mu T_9})^{3/2} \frac{G_{1}G_{2}}{G_{3} (1 + \delta_{12})} e^{-11.605Q/T_9}~,\nonumber \end{eqnarray} where now $N_A$ is Avagadro's number so that $\bigl[ N_A \langle \sigma v(T_9) \rangle^* \bigr]$ is in units of cm$^3$ mol$^{-1}$ s$^{-1}$, $Q$ is in units of MeV and $\hat \mu$ is the reduced mass in atomic mass units. Equations (\ref{eq:thermal}) and (\ref{eq:thermalT9}) are in a convenient form because in the limit of $R \rightarrow 0$, they reduce to usual photodisintegration rates available from various compilations (e.g., \cite{fow67,fow75}; \cite{Holmes76}; \cite{Woosley78}; \cite{CF88}; NACRE, \cite{NACRE}; TALYS, \cite{Goriely08}; NONSMOKER, \cite{Rauscher00} or REACLIB, \cite{REACLIB}). The combined factors multiplying $ {\langle\sigma v\rangle_{12}^*} $ in Eq.~(\ref{eq:thermalT9}) are usually referred to as the "reverse ratio" as this factor gives the reverse reaction rate in terms of the forward rate. In this work, we show that there is a simple correction $(1 + R)$ to this reverse ratio due to the difference between Planckian and Maxwell-Boltzmann statistics. For most of the remainder of this manuscript, our goal will be to derive a simple analytic form for $R$ for ease in correcting existing tabularized reverse reaction rates. We will also derive simple analytic approximations to clarify the essential physics of this correction and summarize examples of which astrophysical conditions may be most affected by these correction factors. \subsection{Nuclear Statistical Equilibrium} Before leaving this discussion, however, it is worth emphasizing again that the above rate does not imply equilibrium, but only detailed balance and a thermal population of photons and nuclear excited states. Nevertheless, the situation of equilibrium between capture and photodissociation frequently occurs in astrophysical environments and is referred to as nuclear statistical equilibrium (NSE). It is of note that the conditions of NSE are also modified from the usual Saha equation by the above quantum corrections. Moreover, in conditions of NSE, one sometimes synthesizes nuclei for which $Q/kT \sim 1$ and the corrections can become larger. Examples of this include the formation of nuclei near the proton drip line in the hot hydrogen burning $rp$-process, or the synthesis of nuclei near the neutron drip line in the neutron-capture $r$-process, as discussed below. To see the revised conditions of NSE consider the evolution of a nucleus undergoing rapid particle captures and photodissociation. This can be written as \begin{equation} \frac{dn_1}{dt} = -n_1 n_2 {\langle\sigma v\rangle^*_{12}} + n_3 n_\gamma \langle\sigma c\rangle^*_{\gamma3}~~. \end{equation} The equilibrium condition $(dn_1/dt) = 0$ therefore demands that \begin{equation} \frac{n_1 n_2}{n_3} = \frac{n_\gamma \langle\sigma c\rangle^*_{\gamma3}} {\langle\sigma v\rangle^*_{12}} = \bigl(1+R\bigr) \biggl(\frac{2 \pi \mu kT}{ h^2}\biggr)^{3/2} \frac{G_{1}G_{2}}{G_{3} (1 + \delta_{12})} e^{-Q/kT}~, \label{sahaeq} \end{equation} or in terms of mass fractions and temperature, it is more convenient for stellar models to write \begin{equation} \frac{X_1 X_2}{X_3} = 9.8685 \times 10^9 \bigl(1+R\bigr) \frac{T_9^{3/2} \hat \mu^{5/2}}{\rho} \frac{G_{1}G_{2}}{G_{3} (1 + \delta_{12})} e^{-11.605Q/T_9}~. \label{sahaeqT9} \end{equation} In the limit that $R \rightarrow 0$, Eqs.~(\ref{sahaeq}) and (\ref{sahaeqT9}) represent the usual nuclear Saha equation \citep{saha21} of statistical equilibrium which also invokes the Maxwellian approximation given in Eq.~(\ref{approx}) either directly or indirectly in its derivation (cf.~\cite{cla68,Iliadis07}). The deviation of NSE due to quantum statistics may impact the evolution of explosive nucleosynthesis environments for which one can encounter nuclei with small photodissociation thresholds, e.g., near the neutron or proton drip lines. To the extent that such nuclei are beta-decay waiting points, for example, the altered statistics will affect the timescale for the build up of abundances. Another possible application of the corrections deduced here is for the ionization equilibrium of atomic or molecular species with a low ionization potential in stellar atmospheres. However, we will not consider that case further here.

Since photons are massless, they are sensitive to the effects of quantum statistics. As a result, at high temperature some of the reverse reaction rates with low $Q$-values can be modified from the tabulated values based upon the approximation of classical Maxwell-Boltzmann statistics for the photons. Moreover, in any environment the quantum effects always speed up the photodisintegration rate because a Planck distribution places more photons at low energies than a Maxwell-Boltzmann distribution of the same temperature. As a result, those nuclei with loosely bound nucleons lose them faster than predicted by the classical Maxwell-Boltzmann distribution. In this paper, we have derived analytic correction terms for the effects of the quantum statistical distribution of photons on tabulated thermonuclear photodisintegration rates. Usually, this modification is small because $(Q/kT) \gg 1$ so that there is little difference between a Planck and a Maxwell-Boltzmann distribution. The effect is largest for environments for which synthesized nuclei have $Q/kT \sim 1$, however, the correction is still small compared to the uncertainty in the estimated thermonuclear reaction rates for such nuclei. We have analyzed possible effects of these corrections in a variety of astrophysical environments including the neutron-capture $r$-process, the hot hydrogen burning $rp$-process, core or explosive silicon burning, the photonuclear $\gamma$-process and big bang nucleosynthesis. In general these corrections have little effect except, perhaps in the case of the $rp$-process for those reactions near the proton drip line waiting points, or in the early stages of the $r$-process when the neutron density is high enough to drive the $r$-process path to nuclei with low $Q$-values even at high temperature.

1012.0519

1012

1012.0343_arXiv.txt

Studies of element abundances in stars are of fundamental interest for their impact in a wide astrophysical context, from our understanding of galactic chemistry and its evolution, to their effect on models of stellar interiors, to the influence of the composition of material in young stellar environments on the planet formation process. We review recent results of studies of abundance properties of X-ray emitting plasmas in stars, ranging from the corona of the Sun and other solar-like stars, to pre-main sequence low-mass stars, and to early-type stars. We discuss the status of our understanding of abundance patterns in stellar X-ray plasmas, and recent advances made possible by accurate diagnostics now accessible thanks to the high resolution X-ray spectroscopy with \cha\ and \xmm.

\label{intro} The determination of the chemical composition of plasmas is of fundamental importance in very different areas of astrophysics. The element abundances in stellar atmospheres have significant impact on the enrichment of the interstellar medium, the evolution of stellar galactic populations, star formation processes, and the structure of stellar interiors. The solar chemical composition provides the standard reference for the elemental abundances studies of other astronomical objects \citep[see review by][]{Asplund09}. However, the composition of solar plasmas is not uniform, and in the outer atmosphere is not constant. Evidence for abundance anomalies in the solar corona with respect to the solar photospheric composition arose from early spectroscopic studies of the solar upper atmosphere. Indeed, the solar corona possesses a chemical composition that is similar to that of the solar wind and solar energetic particles, and at variance with the underlying photosphere \citep[see reviews by][]{Meyer85,Feldman92}. These abundance anomalies, which will be briefly reviewed in \S\ref{ssec:solar}, reflect the effect of still unknown physical mechanisms of chemical fractionation in the process of mass transport into the corona. X-ray spectra of other stars provide us with a means to investigate whether, similarly to the Sun, the chemical composition of the stellar outer atmospheres is different from their underlying photospheric mixture. Furthermore, stellar studies allow us to study the fractionation processes as a function of a wide range of stellar parameters, not accessible from solar studies alone. In this paper, we review the current understanding of the chemical composition of the X-ray emitting plasma in stars, as derived from stellar observations in the EUV and X-ray bands during the past two decades, focusing on recent results from high-resolution X-ray spectroscopy. X-ray observations at high spectral resolution now available with \cha\ and \xmm\ allow us to better disentangle the temperature and abundance effects on the stellar X-ray emission, and represent a significant advance in building a robust scenario for the abundance properties of stellar outer atmospheres. In \S\ref{sec:coronae} we present a review of abundances studies for the solar corona (\S\ref{ssec:solar}) and the coronae of other stars (\S\ref{ssec:stars}). A short discussion of some theoretical studies attempting the modeling of these observed features is included in \S\ref{ssec:models}. In \S\ref{sec:massive} an overview of abundance studies in X-ray spectra of early-type stars is presented, and finally in \S\ref{sec:tts} we review studies of abundances in pre-main sequence low-mass stars and the possible insights they offer into the physical processes producing X-rays in these young stars.

\label{sec:conclusions} In this paper we have reviewed recent advances in our understanding of the chemical composition of X-ray emitting plasma in stars brought about in the past decade. In particular, high-resolution X-ray spectroscopy with \cha\ and \xmm\ has allowed robust determination of the element abundances from X-ray spectra indicating that the abundance anomalies in stellar coronae are real and not an artifact of the modeling of low and medium resolution spectra. Abundance anomalies in coronae of cool stars are largely described by fractionation processes dependent on the element's first ionization potential, and they appear to be a function of the stellar X-ray activity level. A solar-like FIP effect (abundance enhancement of low-FIP elements in coronal plasma) is typically observed in other low to intermediate activity stars similar to the Sun, whereas high activity stars are characterized by an inverse FIP effect (depletion of low-FIP elements in the corona). These findings, however, heavily rely on the assumption that the often unknown underlying stellar photospheric abundances are similar to the solar photospheric abundances. More photospheric abundance studies are needed in order to uncover the true coronal abundance anomalies in a more reliable way and firmly establish the validity of the apparent trends. The abundance of neon shows an interesting pattern with active stars showing a Ne/O abundance ratio significantly larger than the Sun and other low activity stars. This might point to a fractionation of Ne in coronal plasma, either depletion in solar-like activity stars or enhancement in active stars, and raises the issue of what the photospheric Ne abundance is in the Sun and in nearby stars. Studies of flares suggest significant abundance variations compared to quiescent conditions, though the limited quality of the time-resolved spectroscopy achievable at present, and the effects of non-equilibrium, which are difficult to model and constrain, cast some doubts on the robustness of these findings. Improved capabilities for temporally resolved spectral diagnostics are required to be able to confirm these results. Recent studies also suggest that the chemical fractionation of coronal plasma might depend on the stellar spectral type. This dependence, however, appears to apply only to low to intermediate activity stars, with possibly important consequences in the context of developing a theoretical framework to interpret the abundance anomalies. In fact, this finding can put significant constraints on the models as it might suggest that instead of a unique mechanism of fractionation for all coronae, different processes might be dominant in stars similar to the Sun and in stars with higher activity level. For massive stars, an accurate knowledge of the abundances of their X-ray emitting plasmas is still lacking. However, recent investigations with high-resolution spectra have produced interesting results, suggesting subsolar Fe abundance for several early-type stars, and a possible temperature dependent solar-like FIP effect in some sources with hard X-ray spectra. In young low-mass stars the chemical composition of X-ray plasmas typically shows characteristics analogous to more evolved stars with similar activity levels, i.e., an inverse FIP effect, in particular with low Fe abundance and high Ne. Also for pre-main sequence stars more reliable photospheric abundances need to be determined in order to establish the actual extent of these apparent coronal abundance anomalies. The coronal abundances of T Tauri stars do not appear to depend on the presence of ongoing accretion or lack thereof. A few unusually old accreting T Tauri stars show peculiarly high Ne/O abundance ratio in their X-ray emission; this anomalous chemical composition is suggestive of advanced evolution of the circumstellar disk yielding depletion of grain-forming elements, compared to Ne which is volatile. These abundance peculiarities, however, are not found consistently in other old accreting T Tauri stars therefore questioning the validity of this scenario. Although the study of the element abundances in stellar X-ray plasmas has recently yielded significant progress, substantial improvements on both the observational constraints and theoretical models are required to begin understanding the physics of chemical fractionation in stars. This process is likely connected to the yet poorly understood mechanism(s) of mass and energy transport to the corona, which are among the most fundamental open issues in astrophysics.

1012.0343

1012

1012.5085_arXiv.txt

We examine the rest-frame far-infrared emission from powerful radio sources with 1.4\,GHz luminosity densities of $25$$\le$$\log(L_{1.4}$/WHz$^{-1})$$\le$$26.5$ in the extragalactic {\it Spitzer} First Look Survey field. We combine {\it Herschel}/SPIRE flux densities with {\it Spitzer}/IRAC and MIPS infrared data to obtain total ($8-1000\,\mu$m) infrared luminosities for these radio sources. We separate our sources into a moderate, $0.4$$<$$z$$<$$0.9$, and a high, $1.2$$<$$z$$<$$3.0$, redshift sub-sample and we use {\it Spitzer} observations of a $z$$<$$0.1$ 3CRR sample as a local comparison. By comparison to numbers from the SKA Simulated Skies we find that our moderate redshift sample is complete and our high redshift sample is 14\,per cent complete. We constrain the ranges of mean star formation rates (SFRs) to be $3.4$$-$$4.2$, $18$$-$$41$ and $80$$-$$581\,{\rm M}_\odot$yr$^{-1}$ for the local, moderate and high redshift samples respectively. Hence, we observe an increase in the mean SFR with increasing redshift which we can parameterise as $\sim(1+z)^Q$, where $Q=4.2\pm0.8$. However we observe no trends of mean SFR with radio luminosity within the moderate or high redshift bins. We estimate that radio-loud AGN in the high redshift sample contribute $0.1$$-$$0.5\,$per cent to the total SFR density at that epoch. Hence, if all luminous starbursts host radio-loud AGN we infer a radio-loud phase duty cycle of $0.001$$-$$0.005$.

There is now strong evidence that powerful active galactic nuclei (AGN) played a key role in the evolution of galaxies. The correlation of central black hole and stellar bulge mass \citep{Magorrian:98}, and the increased prevalence of star formation \citep{Hopkins:06, Giavalisco:04b} and AGN activity \citep{Wall:05,Aird:10} at earlier epochs suggest that the growth of the black hole is somehow related to the growth of the host galaxy. In the local Universe we see little evidence of high star formation rates (SFRs) in galaxies with powerful radio-loud AGN activity \citep[e.g][]{Condon:98b,Mauch:07}. In the distant Universe, $z>1$, luminous radio galaxies \citep{Seymour:07} and powerful starbursts \citep{Borys:05,Casey:09} are both hosted by massive galaxies, suggesting a common parent population. The idea that these processes are likely to be connected at the epoch when black holes and galaxies went through their most rapid phases of growth has been invoked within various semi-analytical models in order to reconcile those models with observations \citep[e.g.][]{Springel:05}. This connection between central black hole growth and star formation rate is often considered in the context of `feedback' process(es), as the former is postulated to regulate the latter. In particular there is observational evidence, as well as theoretical models, in which the jet from an AGN can produce either positive or negative feedback, where the jet, traced by its radio emission, stimulates or quenches star formation respectively. There is some observational evidence of {\em positive} feedback, whereby star formation is triggered by an AGN jet, e.g. in Minkowski's Object by a jet from NGC 541 \citep{vanBreugel:85,Croft:06}, as well as theoretical models which suggest that the shocks generated by jet propagation can trigger collapse of over-dense clouds and lead to star formation \citep{Fragile:04, Saxton:05}. {\em Negative} feedback by AGN jets would likely require the removal of fuel for star formation, evidence for which are the powerful AGN-induced outflows which have been seen in high redshift radio galaxies \citep{Nesvadba:06,Nesvadba:08}. Such a scenario has also been proposed to regulate the growth of massive galaxies in semi-empirical models \citep{Croton:06, Bower:06}, but this process is only important globally at late times, $z<1$. At earlier times it would be most important in halting the growth of the most massive galaxies. Star formation in powerful AGN has been difficult to trace so far. This difficulty is due to heavy contamination in traditional diagnostics by emission from the AGN (e.g. UV luminosity or optical emission line strengths) as well as obscuration by gas and dust. However the far-infrared (far-IR) presents a window in the electromagnetic spectrum where AGN emission is weak and star formation, if present, can dominate. AGN dust emission tends to peak in the near/mid-IR so far-IR emission should be a cleaner measure of SFR than other traditional methods. It is also possible to use the near/mid-IR to model and subtract any potential AGN contribution to the far-IR \citep[e.g.][]{Hatziminaoglou:08}. There is evidence for extreme SFRs in many powerful high redshift radio galaxies \citep[$z>2$, 1.4\,GHz luminosity densities, $L_{1.4}\ge10^{27}\,$WHz$^{-1}$,][]{Miley:08} from their strong sub-mm emission \citep{Archibald:01, Reuland:04, Greve:06}, their mid-IR spectra \citep[][J. Rawlings, 2011, in prep.]{Seymour:08b} and the spectacular ($>$$100\,$kpc) Ly$\alpha$ haloes sometimes observed \citep{Reuland:03, VillarMartin:03} showing the extended gas that can provide the fuel for star formation. To compliment future targetted {\it Herschel\,} studies of the rare, very powerful radio-loud AGN, we examine in this work less luminous radio-loud AGN, $26.5>\log(L_{1.4}$/WHz$^{-1})\ge 25$, which can be found in reasonable abundance over areas of a few square degrees. We use this definition of `radio-loud' AGN, based on radio luminosity density \citep[\eg][]{Miller:90}, in order to avoid making any distinction between type 1 and type 2 AGN, i.e. AGN classification based upon optical spectroscopy, where different amounts of AGN obscuration may affect the relative amount of optical emission. Although, as we shall show, most of these sources are also `radio-loud' when using the definition of \citet[][5\,GHz over $B-$band luminosity $>10$]{Kellerman:89}. Star formation in these less luminous radio-loud AGN remains poorly studied, as there has been no systematic follow-up of such sources above $z>0.1$. Recently, the importance of radio-loud AGN in this luminosity range was demonstrated by \citet{Sajina:07} who found that $40\,$per cent of $z\sim2$ ULIRGs with deep silicate absorption features were radio-loud and those authors postulated that such sources are transition `feedback' objects after the radio jet has turned on, but before feedback has halted black hole accretion and star formation. The SPIRE instrument \citep{Griffin:10} on board the {\it Herschel Space Observatory} \citep{Pilbratt:10} gives us a clear view of the far-IR/sub-millimeter Universe at wavelengths where many galaxies emit most of their luminosity. The Herschel Multi-tiered Extragalactic Survey (HerMES\footnote{http://hermes.sussex.ac.uk}, Oliver et al. 2011, in prep) provides deep infrared SPIRE data over many of the best studied extra-galactic survey fields. Recent results from {\it Herschel} show that SPIRE detected AGN in deep HerMES fields have far-IR colours similar to the bulk of the SPIRE population which are believed to be star formation dominated \citep{Elbaz:10, Hatziminaoglou:10} and modeling of their spectral energy distributions (SEDs) suggests the SPIRE emission in AGN is dominated by a star forming component \citep{Hatziminaoglou:10}. The work presented here uses {\it Herschel}/SPIRE observations of the {\it Spitzer} Extragalactic First Look Survey (FLS) field taken as part of the {\it Herschel} Science Demonstration Phase (SDP) in October to November 2009. Of the fields observed in SDP this field had the best combination of wide area, uniform radio coverage and good multi-wavelength follow-up. We present our sample of moderate and high redshift radio-loud AGN in \S2 and, we derive IR luminosities and star formation rates in \S3. We present our results in \S4 and discuss them in \S5. We conclude this paper in \S6. Throughout we use a `concordance' cosmology of $\Omega_{\rm M} = 1 - \Omega_{\Lambda} = 0.3$, $\Omega_0 = 1$, and $H_0 = 70\, \kmpspMpc$.

We have examined the incidence of far-IR emission and inferred SFR of luminous radio-loud AGN in a moderate redshift, $0.4<z<0.9$, and a high redshift sub-sample, $1.2<z<3$, as well as a local, $z<0.1$, comparison sample. We have: \begin{itemize} \item constrained the mean SFR of radio-loud AGN to be $3.4$$-$$4.2$, $18$$-$$41$ and $80$$-$$581\,{\rm M}_\odot$yr$^{-1}$ for the local, moderate and high redshift samples respectively, hence, we measure the evolution of the mean SFR to be $\sim(1+z)^{4.2\pm0.8}$, \item observed no strong trends of SFR with radio luminosity in any redshift bin, \item estimated that the host galaxies of radio-loud AGN in the high redshift sub-sample contribute $0.1-0.5\,$per cent to the total SFR density at that epoch and if all LIRGs and ULIRGs have a radio-loud phase we infer a duty cycle of $0.001-0.005$ in such sources. \end{itemize} These results demonstrate that in the distant Universe a considerable amount of star formation is occuring in galaxies hosting a radio-loud AGN, consistent with the frequent evidence for high SFRs in classic high redshift radio galaxies. The mean SFR evolves more quickly than the IR luminosity function implying that some of the star formation is directly related to the radio-loud AGN activity. Both starburst and active nuclear processes have relatively short time-scales so their co-existence in many objects suggests that bursts of star formation and jet activity either are quite common or connected via `feedback'. But is the jet initiating or quenching star formation, or are the processes independent? We cannot answer such questions here, but we shall be able to do so with follow-up of individual sources (to search for outflows of jet-triggered star formation or for mergers triggering both) and with the huge sample that will be provided by the full HerMES data set combined with improved redshift information.

1012.5085

1012

1012.3176_arXiv.txt

We present a sample of edge-on spiral galaxies both of early and late types. The sample consists of 175 galaxies in the $K_s$-filter, 169 galaxies in the $H$-filter and 165 galaxies in the $J$-filter. Bulge and disc decompositions of each galaxy image, taken from the Two Micron All Sky Survey (2MASS), were performed. We discuss several scaling relations for bulges and discs which indicate a tight link between their formation and evolution. We show that galaxies with bulges fitted by the \ser\ index $n < 2$ (pseudobulges) have quite different distributions of their structural parameters than galaxies with $n \geqslant 2$ bulges (classical bulges). First of all, the distribution of the apparent bulge axis ratio $q_\mathrm{b}$ for the subsample with $n < 2$ can be attributed to triaxial, nearly prolate bulges, while $n \geqslant 2$ bulges seem to be oblate spheroids with moderate flattening. Secondly, the Photometric Plane of the sample bulges is not flat and has a prominent curvature towards small values of $n$. Thirdly, despite of the existence of a clear relation between the flattening of stellar discs $h/z_0$ and the relative mass of a spherical component, the distributions over both parameters are quite different for galaxies possesing classical bulges and pseudobulges.

1012.3176

1012

1012.4756_arXiv.txt

{We analysed high-resolution UVES spectra of six stars belonging to the subgiant branch of $\omega$~Centauri, and derived abundance ratios of 19 chemical elements (namely Al, Ba, C, Ca, Co, Cr, Cu, Fe, La, Mg, Mn, N, Na, Ni, Sc, Si, Sr, Ti, and Y). A comparison with previous abundance determinations for red giants provided remarkable agreement and allowed us to identify the sub-populations to which our targets belong. We found that three targets belong to a low-metallicity population at [Fe/H]$\simeq$--2.0~dex, [$\alpha$/Fe]$\simeq$+0.4~dex and [s/Fe]$\simeq$0~dex. Stars with similar characteristics were found in small amounts by past surveys of red giants. We discuss the possibility that they belong to a separate sub-population that we name VMP (very metal-poor, at most 5\% of the total cluster population), which -- in the self-enrichment hypothesis -- is the best-candidate first stellar generation in $\omega$~Cen. Two of the remaining targets belong to the dominant metal-poor population (MP) at [Fe/H]$\simeq$--1.7~dex, and the last one to the metal-intermediate (MInt) one at [Fe/H]$\simeq$--1.2~dex. The existence of the newly defined VMP population could help to understand some puzzling results based on low-resolution spectroscopy (Sollima et al., Villanova et al.) in their age differences determinations, because the metallicity resolution of these studies was probably not enough to detect the VMP population. The VMP could also correspond to some of the additional substructures of the subgiant-branch region found in the latest HST photometry (Bellini et al.). After trying to correlate chemical abundances with substructures in the subgiant branch of $\omega$~Cen, we found that the age difference between the VMP and MP populations should be small (0$\pm$2~Gyr), while the difference between the MP and MInt populations could be slightly larger (2$\pm$2~Gyr). }

\label{sec-intro} Much has been written on $\omega$~Cen, which is probably the most studied cluster in the Milky Way \citep{fvl02}. From the pioneering studies in the 60s to the latest high-quality data and models, more and more details of its multiple stellar populations have come to light, but the picture did not become as clear as expected. Excellent photometries and astrometric catalogues have recently been produced by both ground-based \citep[such as][]{lee99,p00,fvl00,hilker00,hughes00,sollima05a,calamida09,bellini09} and space telescopes \citep[e.g.,][]{ferraro04,bedin04,bellini10}, complemented by high-quality spectroscopic abundance studies both at high \citep[e.g.,][]{norris95,smith00,cunha02,p03,johnson08,johnson09,johnson10,marino10} and low \citep{norris96,suntzeff96,sollima05b,stanford06,villanova07} resolution. Nevertheless, several puzzles still await solutions. One of the open problems concerns the most sensitive region of the colour magnitude diagram (hereafter CMD) to age differences: the subgiant branch (SGB). A large number of photometric and low-resolution spectroscopic studies \citep{hughes00,hilker00,p03,hughes04,hilker04,rey04,ferraro04,stanford06} found age spreads ranging from 2 to 6~Gyr, with a few exceptions and puzzles (see Section~\ref{sec_ages}, for more details). One example of the difficulties encountered in the study of this complex region is posed by the two studies by \citet{sollima05b} and \citet{villanova07}, who used the same ACS dataset and low-resolution spectra of similar quality, but reached opposite conclusions on the total age spread -- and age distribution -- of $\omega$~Cen. Still, even with the best ACS photometries \citep[see, e.g.][]{bellini10}, it is not easy to understand which features of the SGB region correspond to each of the populations that are spectroscopically identified on the red giant branch (RGB), which are known in great detail thanks to the above cited works. This understanding is crucial to solve the relative ages problem in $\omega$~Cen, and to derive the age-metallicity relation, a fundamental ingredient of any model for the formation and evolution of this unique stellar system. \begin{table*} \caption{Observing Logs.} \label{tab_logs} \begin{center} \begin{tabular}{l c c c c c c c c c l} \hline\hline \noalign{\smallskip} ID$_{WFI}$ & R.A. (J2000) & Dec (J2000) & V & (B-V)$_0$ & (V-I)$_0$ & R & t$_{exp}^{(tot)}$ & S/N & $V_r$ & Notes \\ & (deg) & (deg) & (mag) & (mag) & (mag) & ($\lambda/ \delta\lambda$) & (sec) & (@550nm) & (km s$^{-1}$) & \\ \noalign{\smallskip} \hline \noalign{\smallskip} 503358 & 201.960086 & -47.759426 & 17.53 & 0.65 & 0.65 & 45\,000 & 10800 & 50 & 231.0 $\pm$ 0.2 & \rm{Lower SGB} \\ 503951 & 201.870201 & -47.746107 & 17.58 & 0.65 & 0.65 & 45\,000 & 11400 & 50 & 235.8 $\pm$ 0.1 & \rm{Lower SGB} \\ 507109 & 201.848173 & -47.692322 & 17.31 & 0.60 & 0.66 & 45\,000 & 11288 & 50 & 229.7 $\pm$ 0.5 & \rm{Upper SGB} \\ 507633 & 201.813303 & -47.685166 & 17.39 & 0.56 & 0.66 & 45\,000 & 12000 & 45 & 226.7 $\pm$ 0.3 & \rm{Upper SGB} \\ 512115 & 201.846338 & -47.637524 & 17.64 & 0.65 & 0.66 & 45\,000 & 11400 & 45 & 241.4 $\pm$ 0.1 & \rm{Lower SGB} \\ 512938 & 201.804070 & -47.630694 & 17.32 & 0.60 & 0.65 & 45\,000 & 10000 & 50 & 234.6 $\pm$ 0.2 & \rm{Upper SGB} \\ \hline\hline \end{tabular} \end{center} \end{table*} In this paper we present the analysis of a set of UVES high-resolution spectra of six SGB stars, selected from the wide field photometric catalogue by \citet{p00} and \citet{p03}. A preliminary analysis of the same dataset was presented by \citet{p03}. We describe the spectra reductions in Section~\ref{sec-data}; the abundance analysis in Section~\ref{sec-abo}; and the abundance ratio results in Section~\ref{sec-res}. Our main results are discussed in detail in Section~\ref{sec-disc} and are summarized in Section~\ref{sec-concl}. \begin{figure} \centering \includegraphics[width=\columnwidth]{fig_cmds.ps} \caption{Location of the programme stars on the CMD of $\omega$~Cen. The WFI B, B--I photometry of stars in CCD\#5 \citep[from][]{p00} is shown as grey dots. The six UVES targets are marked (black filled circles) with their WFI catalogue numbers.} \label{fig_cmds} \end{figure} \begin{figure} \centering \includegraphics[width=\columnwidth]{fig_area.ps} \caption{Location of the programme stars on the area of $\omega$~Cen. Grey dots mark stars belonging to the WFI photometry by \citet{p00}. Filled circles mark the postion of the six UVES targets.} \label{fig_area} \end{figure}

\label{sec-concl} We analysed UVES high-resolution spectra of six stars on the SGB of $\omega$~Centauri. We compared our results with RGB high-resolution spectroscopy to identify the sub-populations to which our targets belong, and we found remarkable agreement with past abundance determinations. Three of our targets (WFI~507109, 507633, and 512939) have [Fe/H]$\simeq$--2.0~dex, are $\alpha$-enhanced and show no significant s-process enhancement. Two of the remaining targets (WFI~503358 and 503951) belong to the MP population, with [Fe/H]$\simeq$--1.65~dex, $\alpha$-enhanced and with [s/Fe]$\simeq$+0.5~dex. The last target, star WFI~512115, belongs to the MInt2 population, with [Fe/H]=--1.19~dex, $\alpha$-ehanced and with s-process enhancement similar to the MP targets, i.e., slightly lower than what is expected for MInt stars. Our main result (see Section~\ref{sec_first}) is that {\em there exists an additional, metal-poor population (that we name VMP) at [Fe/H]$\simeq$--2.0~dex, which has chemical properties that make it the ideal candidate (remnant of) the primordial population of $\omega$~Cen.} The RGB star ROA~213 \citet{smith00}, star 85007 by \citet{villanova10}, and 25 red giants studied by \citet{johnson10}, have similar chemical composition and could represent the prototype VMP members along the RGB. In particular, the s-process enhancement of the SGB-VMP population is not compatible with the RGB-MP stars, while it is compatible with the quoted RGB-VMP stars. Our conclusion is also supported by previous work on metallicity distributions or RGB and SGB stars and by the exquisite photometry by \citet{bellini10}. We estimate that this VMP population should comprise at most 5\% of the entire stellar content of $\omega$~Centauri, at present. The presence or absence of light element anti-correlations in this population would be a fundamental constraint to the nature of $\omega$~Cen, because anti-correlations are generally exclusively found in globular clusters and never in the field populations of galaxies. From the available literature \citep[mainly][]{johnson09,johnson10,marino10} it appears that (anti-)correlations could be reduced in extent, in VMP stars. Until the presence of (anti)-correlations in VMP stars is excluded by larger data samples, it looks more promising to interpret this as the primordial population of $\omega$ Cen instead of the remnant field population of its putative parent galaxy. The high-precision abundance determinations obtained allowed us to try to shed some light on the SGB morphology relation with the RGB spectroscopically identified sub-populations. We conclude that {\em there appears to be no one-to-one correspondence between the nicely combed substructures of the RGB and SGB. In particular, the MP and MInt populations could either be mixed along the (lower) A, B, and C branches, or be positioned in a not strictly monotonic order, with MInt1 occupying branch A (lower) and MP branch B, just as an example.} As already said, VMP stars should occupy (and possibly dominate) the uppermost SGB branch, or branch A in the \citet{villanova07} terminology. We also found that (see Section~\ref{sec_ages}) {\em the existence of the VMP population could alleviate some of the problems found in previous determinations of relative ages.} In particular, the small metallicity difference between the VMP and MP populations could have escaped previous abundance analyses based on low-resolution spectra. The puzzling result by \citet{villanova07} that the MP population should contain two groups with different ages could indeed be explained by the metallicity difference between VMP and MP. Also, the (too) small age difference found by \citet{sollima05b} between the MP and MInt populations woul become slightly larger when taking into account the existence of the VMP, which should dominate the uppermost SGB envelope. Finally, {\em there should be a small age difference between the VMP and MP populations (0$\pm$2~Gyr), while a slightly larger age difference (2$\pm$2~Gyr) should occur between the VMP and the MInt2 populations.} Althought this latter result is less secure because it relies on one star only, it agrees very well with the majority of past studies \citep[see, e.g.,][]{stanford06}. The use of different sets of isochrones (Padova, BaSTI) does not change the result significantly, and the helium abundance problem should have a negligible impact on the MInt2 star WFI~512115 (the only one which should have higher helium) because, to first approximation, it should change its [Fe/H] by $\simeq$0.08~dex. The age distribution suggested by the present data would accommodate a fast enrichment between the VMP and MP populations, dominated by SNe type II, while the s-process enrichment of the MP ([s/Fe]$\simeq$+0.5~dex) could still pose a problem. The age difference between the MP and MInt populations could instead be sufficient to allow for some intermediate-mass AGB star enrichment \citep{busso99}, bringing [s/Fe] to +1.0~dex. We conclude by noting that this is the only high-resolution-based abundance analysis published on SGB stars in $\omega$~Cen so far. Even if the precision of the abundances is higher than in past low-resolution studies, of course the number of stars examined is only six. To give the final answer to the relative ages problem in $\omega$~Cen, and to identify who is who in the CMD at the SGB level, a much larger sample (a few hundreds) of relatively high-resolution spectra in the SGB region is absolutely necessary.

1012.4756

1012

1012.4426_arXiv.txt

We use the $W_{\Ha}$ versus \nii/\Ha (WHAN) diagram introduced by us in previous work to provide a comprehensive emission-line classification of Sloan Digital Sky Survey galaxies. This classification is able to cope with the large population of weak line galaxies that do not appear in traditional diagrams due to a lack of some of the diagnostic lines. A further advantage of the WHAN diagram is to allow the differentiation between two very distinct classes that overlap in the LINER region of traditional diagnostic diagrams. These are galaxies hosting a weakly active galactic nucleus (wAGN) and ``retired galaxies'' (RGs), i.e. galaxies that have stopped forming stars and are ionized by their \emph{hot evolved low-mass stars} (HOLMES). A useful criterion to distinguish true from fake AGN (i.e. the RGs) is the value of $\xi$, which measures the ratio of the extinction-corrected \Ha\ luminosity with respect to the \Ha\ luminosity expected from photoionization by stellar populations older than $10^{8}$\,yr. We find that $\xi$ follows a markedly bimodal distribution, with a $\xi \gg 1$ population composed by systems undergoing star-formation and/or nuclear activity, and a peak at $\xi \sim 1$ corresponding to the prediction of the RG model. We base our classification scheme not on $\xi$ but on a more readily available and model-independent quantity which provides an excellent observational proxy for $\xi$: the equivalent width of \Ha. Based on the bimodal distribution of $W_{\Ha}$, we set the practical division between wAGN and RGs at $W_{\Ha}=3$ \AA. Five classes of galaxies are identified within the WHAN diagram: \begin{itemize} \item Pure star forming galaxies: $\log \nii/\Ha\ < -0.4$ and $W_{\Ha} > 3$ \AA \item Strong AGN (i.e., Seyferts): $\log \nii/\Ha\ > -0.4$ and $W_{\Ha} > 6$ \AA \item Weak AGN: $\log \nii/\Ha\ > -0.4$ and $W_{\Ha}$ between 3 and 6 \AA \item Retired galaxies (i.e., fake AGN): $W_{\Ha} < 3$ \AA \item Passive galaxies (actually, line-less galaxies): $W_{\Ha}$ and $W_{\nii} < 0.5$ \AA \end{itemize} A comparative analysis of star formation histories and of other physical and observational properties in these different classes of galaxies corroborates our proposed differentiation between RGs and weak AGN in the LINER-like family. This analysis also shows similarities between strong and weak AGN on the one hand, and retired and passive galaxies on the other.

The goal of a classification scheme is to rationally organize a large collection of objects (animals, plants, books, galaxies) into a fixed number of smaller classes \emph{without leaving any object aside}, in order to help dealing with this large collection. There are obviously many ways to classify objects, depending on what properties one is mostly interested in. For galaxies, one could, for example, be mainly interested in the morphology \citep{Hubble1936}, colours \citep{deVaucouleurs1960}, spectral features \citep{Morgan.Mayall1957} etc. It is not only the considered object properties that change from one classification system to another, but also the philosophies underlying the classification process (hierarchical or non hierarchical, supervised or not supervised, \ldots) all of which have their advantages and drawbacks. Ideally, a classification scheme should be objective and allow easy incorporation of new objects into predefined categories. The last criterion, for instance, is not fulfilled by the classifications using principal component analysis \citep{Sodre.Cuevas1994, Connolly.etal1995}. Most importantly, the classification scheme must be useful, Hubble's famous tuning fork being a classical example. Another one is the emission-line classification scheme pioneered by \citet*{Baldwin.etal1981} (see also \citealt{Pastoriza1968}), which has been widely used over the last three decades. The advent of large surveys of galaxies providing extensive collections of homogeneous data sets not easily tractable with conventional methods fostered new classification schemes based on automatic procedures such as neural networks \citep*{Folkes.etal1996}, k-means cluster analysis of entire galaxy spectra \citep{SanchezAlmeida.etal2010}, or on human-eye analysis of galaxy morphology by hundreds of thousands of volunteers, as in the beautiful Galaxy Zoo project \citep{Lintott.etal2010}. Emission-line classifications of galaxies are among the easiest to carry out (once the emission line intensities are measured) and allow one to deal with such issues as star formation, chemical composition or nuclear activity. They have, for example, helped to gain insight into the nature of warm infrared galaxies \citep{deGrijp.etal1992}, of luminous infrared galaxies \citep*{Smith.etal1998}, of galaxies in compact groups \citep{Coziol.etal1998} or into the connection between active galactic nuclei (AGN) and their host galaxies \citep{Kauffmann.etal2003c}. The most famous emission-line diagnostic diagram is that based on the \Oiii/\Hb\ vs \Nii/\Ha\ line ratios\footnote{In the remaining of this paper, \Oiii\ and \Nii\ will be denoted \oiii\ and \nii, respectively.} (dubbed the BPT diagram after \citealt{Baldwin.etal1981}). This diagram was built to pin down the main source of ionization in the spectra of extragalactic objects, and became one of the major tools for the classification and analysis of emission line galaxies in the Sloan Digital Sky Survey (SDSS, \citealt{York.etal2000}). As emphasized by \citet[CF10]{CidFernandes.etal2010}, the BPT diagram leaves unclassified a large proportion of emission line galaxies in the SDSS, due to quality requirements on 4 emission lines. The \citet{Kewley.etal2006} classification is even more demanding, as it requires the intensities of as many as 7 lines. To remedy this situation, \citetalias{CidFernandes.etal2010} have proposed a much more economic diagram that allows one to attain the same objectives using only two lines, \Ha\ and \nii, which are generally the most prominent in the spectra of galaxies. This is a diagram plotting the \Ha\ equivalent width versus \nii/\Ha, which was dubbed the EWH$\alpha$n2 diagram in that paper and which we will from now on refer to as the WHAN diagram, for simplicity \footnote{Our WHAN diagram is, in a way, very similar to the $W_{\nii}$ vs \nii/\Ha\ diagram of \cite{Coziol.etal1998}. However, the WHAN diagram is easier to interpret, since it concerns two physically independent quantities, $W_{\Ha}$ measuring the amount of ionizing photons absorbed by the gas relative to the stellar mass, while \nii/\Ha is a function of the nitrogen abundance, of the ionization state and temperature of the gas.}. The border lines between different classes of galaxies in the WHAN diagram were defined by optimal transpositions of the \citet{Kauffmann.etal2003c}, \citet[S06]{Stasinska.etal2006} and \citet{Kewley.etal2006} border lines onto the $W_{\Ha}$ vs \nii/\Ha\ plane. The categories of emission line galaxies considered by \citet{Kauffmann.etal2003c} and \citet{Kewley.etal2006} comprised star forming galaxies (SF), Seyferts, and LINERs\footnote{In those papers, LINERs should have rather been called ``galaxies with LINER-like spectra'', given that LINER is the abbreviation of "low ionization {\em nuclear} emission regions" (as originally introduced by \citealt{Heckman1980}), and the SDSS spectra cover much more than the nuclear regions of galaxies.} This latter category is supposed to comprise objects hosting low-level nuclear activity. However, as demonstrated by \citet[S08]{Stasinska.etal2008}, the LINER region in the BPT diagram also contains galaxies that have stopped forming stars and are actually ionized by the hot evolved low-mass stars (HOLMES) contained in them. This idea is not new. It was already put forward to explain the emission line spectrum of elliptical galaxies \citep{Binette.etal1994, Macchetto.etal1996} and even earlier in different contexts \citep{Hills1972, Terzian1974, Lyon1975, Sokolowski.Bland-Hawthorn1991}. This idea is recently gaining support \citep{Schawinski.etal2010, Masters.etal2010, Kaviraj2010}, especially with the detailed study of nearby galaxies \citep*{Sarzi.etal2010, Annibali.etal2010, Eracleous.etal2010}. As shown by \citetalias{CidFernandes.etal2010}, when taking into account the ``forgotten population'' of weak line galaxies, the proportion of such ``retired galaxies'' (RGs) increases dramatically. It is therefore essential to find ways to distinguish these RGs from truly ``active'' galaxies, i.e. galaxies which contain a (weak) AGN. This is one of the goals of the present paper. Our second goal is to link the universe of emission line galaxies (ELGs) with the universe of galaxies without emission lines (often dubbed ``Passive Galaxies'', PG). We will see that RGs and passive galaxies are in fact very similar objects. With our WHAN diagram, we will thus be able to provide a comprehensive classification of all SDSS galaxies, including PGs. The paper is structured as follows: Section \ref{sec:Data} presents the samples of galaxies used to work out our classification and our {\sc starlight} data base from which the properties of galaxies that we use are extracted. Section \ref{sec:preliminary} lists the initial definition of galaxy classes: SF, Seyferts, LINERs and PGs. Section \ref{sec:RGs}, the core of this paper, sorts out RGs from AGN among LINER-like galaxies and proposes a practical criterion to identify them. Section \ref{sec:Revised_EWHan2} describes our comprehensive classification of galaxies based on the WHAN diagram and leading to five classes: SF, strong AGN, weak AGN, RG and PG. Section \ref{sec:SFHs} examines the star formation histories of galaxies within our new classification scheme, placing RGs in the context of the galaxy population in general. Section \ref{sec: distributions} discusses the distributions of a series of physical and observed properties in the different galaxy classes, to test the pertinence of our new classification. Finally, Section \ref{sec:summary} summarizes our main results.

\label{sec:summary} The goal of this paper was to provide a comprehensive classification of galaxies according to their emission line properties which would be able to cope with the large population of weak line galaxies that do not appear in traditional diagrams (e.g., BPT) because they lack some of the diagnostic lines. This is possible using the WHAN diagram ($W_{\Ha}$ vs \nii/\Ha) that we proposed in earlier work \citepalias{CidFernandes.etal2010}. An additional problem with the BPT diagram and other similar line-ratio diagrams is that, as shown by \citet{Stasinska.etal2008}, the LINER region contains a mix of two completely different families of galaxies: galaxies that contain a weak AGN and ``retired galaxies'', i.e. galaxies that have stopped forming stars and whose emission lines are powered by their hot evolved low-mass stars (HOLMES). A truly comprehensive classification scheme must account for this phenomenon, and the WHAN diagram allows so. To distinguish true AGN from fake ones (i.e., the RGs), we have shown that a useful criterion is the value of $\xi$, which measures the intrinsic extinction-corrected \Ha\ luminosity in units of the \Ha\ luminosity expected from stellar populations older than $10^{8}$\,yr uncovered by our stellar population analysis software {\sc starlight}. In principle, $\xi$ should be equal to one if the galaxy absorbs all the ionizing photons provided by its HOLMES, and $> 1$ if any other source of ionization is present. We have shown that this index is bimodally distributed, with a $\xi \gg 1$ population composed by galaxies undergoing vigorous star-formation and/or nuclear activity, and a population of galaxies centred right at the predicted value for RGs. Computing $\xi$ for each galaxy, however, is not a trivial task: it requires a stellar population analysis able to identify the populations of HOLMES and a recipe to obtain the ionizing radiation field from these HOLMES. The latter strongly depends on how the (yet quite uncertain) evolutionary tracks for post-asymptotic giant branch stars are incorporated into the stellar population codes as well as on other issues such as the IMF and the metallicities of the stellar populations. It is therefore preferable to base the distinction between true and fake AGN on the \Ha\ equivalent width, an easy to obtain observational parameter which provides an excellent proxy for $\xi$. In view of the strong dichotomy in the observed $W_{\Ha}$ distribution for galaxies in the right wing of the BPT diagram (i.e., the wing previously supposed to contain only AGN), we propose a practical separation between weak AGN and RGs based on a class frontier at $W_{\Ha}=3$ \AA. This physically inspired but data-guided separation corresponds to the minimum between the two modes in the $W_{\Ha}$ distribution. We thus identify 5 classes of galaxies in the WHAN diagram: \begin{itemize} \item SF: Pure star forming galaxies: log \nii/\Ha\ $< -0.4$ and $W_{\Ha} > 3$ \AA \item sAGN: Strong AGN (i.e., Seyferts): log \nii/\Ha\ $> -0.4$ and $W_{\Ha} > 6$ \AA \item wAGN: Weak AGN: log \nii/\Ha\ $> -0.4$ and $3 <W_{\Ha} < 6$ \AA \item RG: Retired galaxies (i.e., fake AGN): $W_{\Ha} < 3$ \AA \item PG: Passive galaxies (actually, line-less galaxies): $W_{\Ha}$ and $W_{\nii} < 0.5$ \AA \end{itemize} Note that the border between RGs and PGs is somewhat arbitrary, but this is not really a concern. We have then analysed the star formation histories of the different classes of galaxies, as well as the distribution of a dozen of physical or observational properties, such as the galaxy stellar mass $M_\star$, surface densities $\mu_\star$, velocity dispersions $\sigma_\star$, and stellar extinctions $A_V^\star$ obtained through {\sc starlight}. All these comparisons corroborate our proposed differentiation between RGs and wAGN in the LINER-like family. We also find that wAGN have properties overlapping with those of sAGN, while RGs line up with PGs. In other words, wAGN are the low activity end of the AGN family, while RGs are passive galaxies with emission lines. As a matter of fact, both RGs and PGs are ``retired'' or ``passive'' galaxies, the real spectroscopic difference between them being the presence or absence of emission lines. It is only the inherited custom which imposed the adopted nomenclature. Now that these emission line classes have been established, further work could be to understand what is at the root of the strength of the AGN phenomenon or why some passive galaxies present emission lines and others do not.

1012.4426

1012

1012.4610_arXiv.txt

We have studied the very long-term temporal properties of the optical emission from Be X-ray binaries (BeX) in the Small Magellanic Cloud over a $\sim$ 16 yr baseline, using light curves from the MACHO and OGLE databases. All the BeX in our sample display superorbital variations, many of them quasi-periodic on timescales of $\sim$ 200-3000 d. These long-term variations are believed to be related to the formation and depletion of the circumstellar disc around the Be star and we compare and contrast their behaviour with that of the LMC's prototypical BeX, A0538-66. The great majority of sources show a correlation of outburst amplitude with brightness (the opposite to that seen in A0538-66) although the amplitudes are mostly small ($\le$ 0.1 mag). We suggest this is an orbital inclination effect. In addition, we have also detected many of their optical orbital periodicities, visible as a series of precisely regular outbursts. Furthermore, the amplitude of these periodic outbursts can vary through the long-term superorbital cycle, and we discuss mechanisms which can produce this effect, as well as examining an apparent correlation between these periodicities. As a by-product of this variation survey we have compiled a list of all the reported SMC BeX orbital and superorbital periodicities at optical and X-ray wavelengths.

\label{intro} In high-mass X-ray binaries (HMXBs), a compact object (usually a neutron star) accretes mass from a massive early type O-B star. Conventionally, they are subdivided into two groups, the supergiant X-ray binaries (SgXRB) and Be/X-ray binaries (BeX). These massive X-ray binaries are particularly numerous in the Small Magellanic Cloud (SMC) \citep[e.g.][]{coe05}. From extrapolation of the Milky Way's population of 65 HMXBS and based on the Milky Way/SMC mass ratio being $\approx$ 50, one would expect to find only one or two HMXBs in the SMC. Remarkably, 59 of these systems have now been detected, with another surprising result being that all of them are BeX systems, with only one exception, the supergiant system SMC X-1 \citep{coe05,mcgowan07}. A BeX consists of a neutron star orbiting a Be star in a wide (period $\sim$ months) and eccentric orbit. For reasons not fully understood, Be stars rotate very rapidly, leading to outflows from the Be star that form an equatorial disc around it. If sufficiently extended, the neutron star can penetrate this equatorial disc during each periastron passage, giving rise to periodic outbursts that can be observed over a wide range of wavelengths (optical, X-ray,...) \citep[e.g.][]{okneg01}. A review of basic properties of BeX can be found in \citet{neg98} and \citet{coe05}. Several authors have investigated the extensive optical light curves of these sources available from the MACHO and OGLE projects in search of evidence for orbital modulations \citep{alcock96,udalski97}. In addition to the orbital variations usually seen as a series of periodic outbursts, these sources can show long-term superorbital modulations with timescales of hundreds of days to years. The prototype for this behaviour is the highly luminous ($\ge$ 10$^{39}$ ergs$^{-1}$ at its peak) LMC BeX, A0538-66, which has a remarkably stable superorbital modulation of 420.82 days, compared to its well-established orbital period of 16.65 days \citep[hereafter MC03]{alcock01,mcgowan03}. It was suggested by MC03 that these long-term modulations were a result of the formation and depletion of the circumstellar disc around the donor star which gives rise to the Be phenomenon. Given the known long-term variations in the Be systems, this discovery raised the question of whether such behaviour was ubiquitous, or somehow related to the presence of the X-ray source. Consequently, in this paper, we investigate this superorbital behaviour in much greater detail within the now substantial BeX population by exploiting the MACHO and OGLE project databases. Together these provide a 16 year baseline for LMC/SMC studies of X-ray binary counterparts. Table~\ref{tab:listBe} gives the position (RA and Dec) as well as alternative name of all known SMC X-ray pulsars. We use the naming convention of \citet{coe05}, where SXP`x' is the SMC X-ray Pulsar with a `x' s pulse period. Table~\ref{tab:listcounterpart} details the MACHO and OGLE counterparts of these SMC X-ray pulsars, their OGLE positions (RA and Dec), all previously reported X-ray and optical orbital periods, and their X-ray luminosities.

In this paper we have exploited the $\sim$16yr timebase of the MACHO and OGLE archives in order to investigate the long-term variability properties of almost all Be/X-ray pulsars currently known in the SMC. Virtually all of them show superorbital variations on timescales of hundreds to thousands of days. While such variations had been known historically for some of the bright, galactic Be systems, this is the first systematic study on these timescales of the BeX SMC population, and it is clear that such variability is extremely common, much of it appearing quasi-periodic in nature. Optical light curves of these BeX systems show a variety of periodicities from as short as $\sim$hours (likely non-radial pulsation of the Be star) to tens of days (likely orbital modulation resulting from the neutron star's regular interaction with mass lost from the Be donor in its eccentric orbit) to many hundreds of days, which we presume to be a property of the circumstellar disc surrounding the Be donor. This latter timescale is seen as both a low and high amplitude modulation, explanations for which range from an oscillation within the disc, to its formation and depletion. We will discuss these properties in more detail. The list of all superorbital and orbital periods found in this report as well as the previously reported X-ray and optical periods are combined in table~\ref{tab:listres}. \subsection{Superorbital periods in BeX systems}\label{superorbital} Long-term non-orbital variations have previously been seen in some BeX systems, such as the 421 d quasi-periodic modulation in A0538-66 \citep{alcock01}. This modulation was the driving force behind the current paper, as it suggested as an observable effect of variations in size of the circumstellar disc around the Be star. \citet{alcock01} reported the additional remarkable property that, in A0538-66, the strength of the 16.5 d orbital outbursts depend very strongly on the brightness of the source and hence are also a function of the 421 d cycle. They occur during optical minimum, but not when the star is at its maximum brightness. For this system, the decline in brightness has been suggested to be related to the formation of an extended, cool, equatorial disc about the Be star, which will mask part of the Be star itself. A0538-66 becomes redder at minimum and bluer at optical maximum, which is what might be expected for a high inclination system. In contrast, we find here that only a few of the SMC BeX sources exhibit this type of behaviour, such as SXP25.5 (see figure~\ref{colormag}). However, for almost all of the SMC sources examined here, the source gets {\it redder} when it brightens (figure~\ref{colormag}). For these systems, we interpret the red continuum as arising from the Be circumstellar disc seen at moderate or low inclination, i.e. the disc adds cooler light to that of the B star. BeX systems can exhibit this type of behaviour if the inclination of the Be equatorial disc is $ \leq90^{o}-\alpha $, where $\alpha$ is the opening angle of the disc (and is likely very small, $\leq$ few degrees). Assuming that all the BeX binaries are oriented at random, the probability of a system having an inclination $i\leq90^{o}-\alpha$ is simply $1 - cos(90^{o}-\alpha)$. If we assume that the opening angle of the disc for a typical BeX source is $\alpha\sim10^{o}$, then the probability of a BeX having an inclination $<80^{o}$ is 0.83. i.e. we would therefore expect the BeX systems to exhibit this type of behaviour rather than that of A0538-66. In our sample of 31 sources with MACHO counterparts we would then expect only 4 or 5 sources to exhibit the ``A0538-type'' of behaviour, and this is consistent with the number (3 sources) that become redder as they get fainter (see figure~\ref{colormag}). In addition, we have seen a correlation between the amplitude of outburst and brightness in some sources (such as SXP7.78, SXP293, SXP755, ...). For those systems, the amplitude of outburst varies as a function of the brightness of the source, the amplitude growing with the source brightness (see figure~\ref{outburstvariation}). This suggests that the size of the Be equatorial disc influences the scale of the interaction between the neutron star and the Be equatorial disk. However, if the Be equatorial disc becomes sufficiently large, then it can influence the entire orbit of the neutron star, producing a type II outburst. In this case, the Be circumstellar disk almost reaches its maximum size because it is then truncated by the neutron star. At this point the neutron star acts to increase the density of the disc rather than extend its size, hence the optical modulation is greatly reduced. This would explain the strange behaviour of SXP7.92. In Figure~\ref{colormag}, the MACHO colour variation of SXP7.92 increases with the brightness of the source. However, its outburst amplitude becomes very weak as it brightens (Figure~\ref{outburstvariation}). For this system, the equatorial disc has become very large (it increases in brightness by $\sim$1 mag) and likely encompasses the entire orbit of the neutron star. For highly variable sources (such as SXP6.85, SXP8.88, SXP15.3, SXP18.3,...), the MACHO colour variations follow the variation in brightness of the source. This confirms that the long-term variation in the light curve is related to the behaviour of the Be circumstellar disc. \subsubsection{Contribution of the extended disc.} Let us consider one of these highly variable SMC BeX sources. Assuming that SXP6.85 is viewed face-on ({\it i$\sim$0$^{\circ}$}) and it has no disc during optical minimum, then the 0.7 mag global change in the light curve represents the brightness contribution from the Be disc. This implies a disc luminosity of $\sim 3.3 \times 10^{35}~erg.s^{-1}$, which corresponds to a total increase in brightness by $\sim$95\% from the optical minimum ($L_{min}\sim 3.5 \times 10^{35}~erg.s^{-1}$). However, for a very high inclination system such as A0538-66 \citep[$i\geq 74^{\circ}.9\pm6^{\circ}.5$, ][]{mcgowan03}, the quiescent state (no disc, only the naked B star) corresponds to the maximum optical brightness, and then when the equatorial disc grows it masks part of the Be star and reduces its optical brightness. The change in brightness in the light curve of A0538-66 is about 0.5 mag. However, the size of the equatorial disc of this source cannot be large because of its short orbital period ($\mathrm {P_{orb}}$=16.65 d, \citealt{mcgowan03}), as the disc will be truncated at a smaller radius. The change in luminosity between no disc and maximum size of disc is about $2.7 \times 10^{35}~erg.s^{-1}$, this corresponds to a decrease in the total brightness of $\sim$37\% from its quiescent state. \subsubsection{Interaction with the neutron star.} It is interesting to note that, in spite of their long orbital periods, the presumed high orbital eccentricity leads to high neutron star velocities during periastron passage. Consider a typical BeX source consisting of a 1.4 $M_{\odot}$ compact object orbiting a 10 $M_{\odot}$ Be star with an eccentricity of {\it e = 0.7} and orbital periods of $P_{orb}$= 17 d and 300 d. At periastron (separation= $1.3$ and $8.9 \times 10^{7}~km$), the neutron star would have a velocity of 443 and 170 $km.s^{-1}$ respectively. The projected equatorial rotational velocity of a Be star ($v~sin~i$) can be estimated from the broadening of its spectral lines, with a mean value of about $250~km.s^{-1}$ \citep{slettebak82}, which is comparable to the neutron star velocity at periastron. Therefore, at periastron, the primary star is rotating close to corotation with the orbiting neutron star. \subsubsection{One-armed oscillation in the equatorial disc.} Be stars are believed to exibit a cyclic variation in their emission line profiles (known as the `V/R' variations; \citealt{McLaughlin61}). It is widely accepted that these variations are caused by global one-armed oscillations in the equatorial discs of Be stars \citep{Kato83, Okazaki91}. The period of this long-term V/R variation is typically in the range of 2-15 years for isolated Be stars \citep[eg.][]{Papaloizou92}. We note that these timescales are very similar to those seen in our light curves (see Table~\ref{tab:listres}). This suggests that the long-term variation in Be star emission line profiles (the `V/R' variation) may be related to the long-term variations we see in the optical light curves. Unfortunately, there are not yet any long-term spectrocopic studies of these SMC BeX sources capable of investigating the correlation between these two (photometric and spectrocopic) long-term variations. The peaks and dips seen in SXP2.37 may possibly be explained as a direct effect of such a precessing elongated disk. For SXP2.37, the outbursts are seen during optical maximum, and then change into dips as the source fades. If we assume that SXP2.37 is a high inclination system, the source is brighter when the elongated disc extends to each side of the Be star, in this case the interaction between the neutron star and the disc can be seen as a maximum. On the other hand, when the elongated disc passes in front it will obscure the Be star, reducing its brightness. In this case, the neutron star passage will enlarge the disc which blocks some light from the hotter Be star, which we see as dips. \subsection{Orbital and Super-orbital period correlation} Apart from the very long-term ($\sim$ hundreds of days) variations that are of primary interest here, we have also seen optical orbital modulations in the light curves of these SMC BeX systems, visible as a series of regular outbursts. These outbursts are interpreted as the periastron passage of the neutron star where it interacts with the Be equatorial disc. At periastron, the circumstellar disc can be perturbed from its stable resonant state, this perturbation will slightly increase its surface area and, hence, its optical brightness \citep{okneg01}. The optical light curves folded on the orbital period have an asymmetric profile, with a faster rise and slower decline. This behaviour has been seen in other BeX systems \citep{alcock01}. Furthermore, in some sources such as SXP46.6, SXP327, SXP348, etc., the optical outburst is not just a single maximum, but can show two peaks every binary cycle, which is seen in the light curves when folded on the orbital period. \citet{mcgowan07}, and \citet{coe09} suggested this behaviour as arising from the misalignment between the spin axis of the Be equatorial disc and the orbital plane of the binary system. Therefore, the neutron star interacts twice with the Be equatorial disc every binary cycle. We note that for SXP18.3, we find an optical periodicity (P= 28.5 d) that is distinct from its presumed orbital period (P= 17.92 d), which is clearly seen only in the MACHO, OGLE II, and first year of the OGLE III light curves. The mechanism that produces this 28.5 d optical modulation is not understood but it is clearly visible in the light curve. We note that it appears especially during the optical minimum. The observed misalignment can be caused by the supernova kick received by the system when the neutron star was born. \citet{brandt95} suggested that an asymmetric supernova explosion can give very large kicks to the newly formed neutron stars which can either disrupt the system if the kick is too strong or lead to a large eccentricity and misalignment between the old and new orbits if the velocity of the kick is smaller. \citet{martin09} suggested that a velocity kick of 265 $km.s^{-1}$ is consistent with the observed misalignments in BeX systems, but too high for the observed eccentricities. In Figure~\ref{PorbPsup}, we have plotted the observed long-term superorbital period of the SMC BeX systems against their orbital periods. It appears that the two periods correlate. We have computed the linear Pearson correlation coefficient to be 0.73 with a p-value of 0.00039 (significant at the 99.9\% confidence level). Some authors have already suggested that these long-term variations might be related to the orbital period \citep{reig05,coe05}, based on the H$\alpha$ EW - P$_{orb}$ relationship of \citet{reig97}, which is in good agreement with the disc truncation model of \citet{okneg01}. If the Be equatorial disc is truncated by the neutron star orbit, a source with a shorter period (large number of periastron passages) would have a smaller equatorial disc than a source with a wider and more eccentric orbit. \begin{table*} \centering \caption{Periodicities found in SMC BeX sources.} \vspace{0.5cm} \scriptsize{ \setlength{\columnsep}{-10pt} \begin{tabular}{lrcrccr}\hline\hline \\ \multicolumn{1}{c}{\(\bf {Short}\)}& \multicolumn{2}{c}{\(\bf {P_{sup}^{\star}}\)}& \multicolumn{2}{c}{\(\bf {P_{orb}}\)}& \multicolumn{2}{c}{\(\bf {Previously~reported }\)}\\ \multicolumn{7}{c}{}\\ \multicolumn{1}{c}{\(\bf {ID} \)}& \multicolumn{1}{c}{\(\bf {Period} \)}& \multicolumn{1}{c}{\(\bf {Semi-amplitude} \)}& \multicolumn{1}{c}{\(\bf {Period(error)} \)}& \multicolumn{1}{c}{\(\bf {Semi-amplitude} \)}& \multicolumn{1}{r}{\(\bf {P_{X-ray^{\dagger}}} \)}& \multicolumn{1}{r}{\(\bf {P_{opt}[ref]} \)}\\ \multicolumn{1}{c}{}& \multicolumn{1}{c}{[d]}& \multicolumn{1}{c}{[mmag]}& \multicolumn{1}{c}{[d]}& \multicolumn{1}{c}{[mmag]}& \multicolumn{1}{c}{[d]}& \multicolumn{1}{c}{[d]}\\ \\ \hline \\ SXP0.09 & [247] & $<$5 & ... & $<$5 & ... & ... \\ SXP0.92 & 2654$\pm$298 & 6 & ... & $<$5 & ... & 51[1] \\ SXP2.37 & ... & $<$5 & 18.58(1)$^{\star \star}$ & 16 & ... & ... \\ SXP2.76 & 2800$\pm$700 & 73 & 82.37(7) & 11 & ... & 82.1[2] \\ SXP3.34 & [495] & $<$5 & 11.09(1) & 14 & ... & ... \\ SXP6.85 & 621$\pm$4 & 201 & 110.0(2) & 26 & 112 & 114[4] \\ SXP7.78 & 1116$\pm$56 & 50 & 44.9(2) & 26 & 44.9 & 44.8[5] \\ SXP7.92 & 397$\pm$2 & 138 & 36.41(2) & 15 & ... & 36.8[14] \\ SXP8.9 & 1786$\pm$32 & 485 & 28.51(1) & 8 & 28.4 & 33.4[6] \\ SXP9.13 & 1886$\pm$35 & 32 & 80.1(1) & 8 & 77.2 & 40.1[8] \\ SXP15.3 & 1515$\pm$23 & 55 & 74.51(5) & 12 & 28 & 75.1[8] \\ SXP18.3 & ... & $<$5 & 17.95(1) & 12 & 17.7 & 17.7[15] \\ SXP22.1 & ... & $<$5 & 75.97(6) & 9 & ... & ... \\ SXP25.5 & ... & $<$5 & 22.50(1) & 16 & ... & ... \\ SXP31.0 & ... & $<$5 & 90.5(1) & 14 & ... & 90.4[2] \\ SXP34.1 & ... & $<$5 & [598] & $<$5 & ... & ... \\ SXP46.6 & ... & $<$5 & 136.4(2) & 9 & 137 & 137[9] \\ SXP59 & ... & $<$5 & 62.10(4) & 8 & 122 & 60.2[10] \\ SXP74.7 & 1220$\pm$64 & 21 & 33.37(1) & 12 & 61.6 & 33.4[11] \\ SXP82.4 & ... & $<$5 & 171(1) & 10 & 362 & ... \\ SXP91.1 & ... & $<$5 & 88.3(1) & 24 & 117 & 88.2[7] \\ SXP101 & 758$\pm$6 & 14 & 21.95(1) & 12 & 25.2 & 21.9[12] \\ SXP138 & 2700$\pm$304 & 172 & [143.1] & $<$5 & 103 & 122[6] \\ SXP140 & 492$\pm$2.4 & 110 & ... & $<$5 & ... & 197[6] \\ SXP172 & ... & $<$5 & 67.88(4) & 22 & 70 & 69.9[6] \\ SXP202A & 1220$\pm$61 & 100 & 71.98(5) & 18 & 91 & ... \\ SXP202B & $\sim$3000 & 108 & 224(1) & 15 & ... & ... \\ SXP264 & $\sim$2000 & 47 & 49.06(2) & 18 & ... & 49.1[10] \\ SXP280 & $\sim$2000 & 88 & 126.4(2) & 15 & 64.8 & 127[2] \\ SXP293 & ... & $<$5 & 59.77(3) & 35 & 151 & 59.7[7] \\ SXP304 & ... & $<$5 & [344] & $<$5 & ... & 520[6] \\ SXP327 & 1274$\pm$143 & 41 & 45.9(2) & 120 & ... & 45.9[16] \\ SXP348 & ... & $<$5 & 94.4(1) & 9 & ... & 93.9[6] \\ SXP455 & 1886$\pm$145 & 110 & 74.96(5) & 21 & ... & 74.7[7] \\ SXP504 & 3448$\pm$119 & 32 & 272(1) & 10 & 265 & 273[10] \\ SXP564 & $\sim$3000 & 69 & 152.4(2) & 15 & 151 & 95.3[7] \\ SXP645 & 2857$\pm$81 & 460 & [135.3] & $<$5 & ... & ... \\ SXP701 & ... & $<$5 & 412(5) & 8 & ... & 412[10] \\ SXP755 & ... & $<$5 & 391(2) & 66 & 389 & 394[7] \\ \\ \hline \multicolumn{7}{p{14cm}}{[1]: \citet{kaspi93}; [2]: \citet{schmidtke06}; [3]: \citet{coe05}; [4]: \citet{mcgowan08}; [5]: \citet{cowley04}; [6]: \citet{schmidtkecow06}; [7]: \citet{schmidtke04}; [8]: \citet{edge05a}; [9]: \citet{mcgowan08}; [10]: \citet{schmidtke05b}; [11]: \citet{schmidtke07b}; [12]: \citet{mcgowan07}; [13]: \citet{edge05b}; [14]: \citet{coe09}; [15]: \citet{schurch09}; [16]: \citet{coe08}}\\ \\ \multicolumn{7}{p{14cm}}{Periods in square brackets are marginally significant (see section 2.3).}\\ \multicolumn{7}{p{14cm}}{$^{\star\star}$ 1 $\sigma$ uncertainty in paratheses (units of last digit). }\\ \multicolumn{7}{p{14cm}}{$^{\star}$Superorbital period.}\\ \multicolumn{7}{p{14cm}}{$^{\dagger}$X-ray orbital period from \citet{galache08}.} \label{tab:listres} \end{tabular} } \end{table*} \begin{figure} \scalebox{0.46}{\includegraphics{Superorb_orb.eps}} \caption{Plot of the BeX superorbital periods found in this work against orbital period. The dashed line represents the best linear fit.} \label{PorbPsup} \end{figure}

1012.4610

1012

1012.5036_arXiv.txt

If Dark Energy introduces an acceleration in the universal expansion then large scale gravitational potential wells should be shrinking, causing a blueshift in the CMB photons that cross such structures (Integrated Sachs-Wolfe effect, [ISW]). Galaxy clusters are known to probe those potential wells. In these objects, CMB photons also experience inverse Compton scattering off the hot electrons of the intra-cluster medium, and this results in a distortion with a characteristic spectral signature of the CMB spectrum (the so-called thermal Sunyaev-Zel'dovich effect, [tSZ]). Since both the ISW and the tSZ effects take place in the same potential wells, they must be spatially correlated. We present how this cross ISW-tSZ signal can be detected in a {CMB-data} contained way {by} using the frequency dependence of the tSZ effect in multi frequency CMB experiments like {\it Planck}, {\em without} requiring the {use} of external large scale structure tracers data. {We find that by masking low redshift clusters, the shot noise level decreases significantly, boosting the signal to noise ratio of the ISW--tSZ cross correlation. We also find that galactic and extragalactic dust residuals must be kept at or below the level of $\sim 0.04$ ($\mu$K)$^2$ at $\ell=10$, a limit that is a factor of a few below {\it Planck}'s expectations for foreground subtraction. If this is achieved, CMB observations of the ISW-tSZ cross correlation should also provide an independent probe for the existence of Dark Energy and the amplitude of density perturbations.}

The primary Cosmic Microwave Background (CMB) and especially its angular power spectrum provides us with powerful constraints on the content of the universe and its evolution. It is now well established that an accurate understanding of the primary CMB power spectrum requires a good understanding of the secondary CMB anisotropies resulting from the interaction of the CMB photons with the matter along the line of sight from the last scattering surface to the observer \citep[see][for a review]{Aghanimrevue08}. The great efforts {undertaken} to understand these secondary anisotropies, in order to best recover the primary CMB, also provide us with powerful independent cosmological probes when the secondary anisotropies are regarded as a source of information rather than contamination. Among those secondary CMB anisotropies, some result from the gravitational interaction of the CMB photons with the potential wells they cross. {One of them is the} Integrated Sachs-Wolfe (ISW) effect, {by which CMB photons experience some blue/redshift as} they pass through large scale time evolving potential wells \citep{SachsWolfe1967}. Since a dark energy like component is expected to affect the growth of large scale structures, {making them shallower}, {a} detection of the ISW effect is an important probe {for} establishing {its} existence - {provided that} the Universe is flat and general relativity is a correct description of gravity - and constraining the equation of state of such a component. Detection claims of the ISW effect arose as soon as the first year WMAP data were released. Those claims were based upon cross correlation analyses of WMAP CMB data and {galaxy density templates built from different surveys}. While most of the first analyses were conducted in real space (i.e., by computing the angular cross-correlation function), subsequently new results based upon Fourier/multipole and wavelet space were presented. The results from the WMAP team on the cross correlation of NRAO Very Large Sky Survey (NVSS) with WMAP data \citep{Nolta2004} were soon followed by other analyses applied not only on NVSS data, but also on X-ray and optical based catalogs like HEA0, SDSS, APM or 2MASS, \citep{Boughn2003-4,Fosalba2003,Scranton2003,Fosalba2004,Afshordi2MASS}. As subsequent data releases from both the CMB and the SDSS side became public, new studies prompted further evidence for significant cross correlation between CMB and Large Scale Structure (LSS) data, e.g., \citet{padmanabhan2005,Cabre2006,Giannantonio2006,Rassat2007}. By that time, wavelet techniques were also applied on NVSS and WMAP data, providing the highest significance ISW detection claims at the level of 3--4 $\sigma$, \citep{Vielva2006,Pietrobon2006,McEwen2007}. The initial effort of \citet{Granett2008}, consisting in stacking voids and superclusters extracted from SDSS data, yielded a very high significance ($\sim 4\sigma$) ISW detection claim. However, it was later found in \citet{Granett2009} that such signal could not be due to ISW only, since a gravitational potential reconstruction from the Luminous Red Galaxy (LRG) sample of SDSS yielded a much lower signal ($\sim 2.5\sigma$). \citet{Giannantonio2008} and \citet{Ho2008} used different LSS surveys in a combined cross-correlation analysis with CMB data and claimed high significance ($\sim 4-5 \sigma$) ISW detections. However, doubts on the validity of such claims have also arised recently. \citet{chm2006} first pointed out the lack of significant cross correlation between WMAP 1st year data and density surveys built upon 2MASS, SDSS and NVSS surveys on the large angular scales, but detected the presence of radio point source emission and thermal Sunyaev-Zel'dovich effect on the small scales. In \citet{chm2008} a study of the expected signal-to-noise ratio (S/N) for different sky coverages was presented, and it was found that in the standard $\Lambda$CDM scenario the ISW -- density cross correlation should be well contained in the largest angular scales, ($l<50-60$). This was proposed as a consistency check for ISW detection against point-source contamination. In \citet{chm2009} cross-correlation analyses between NVSS and WMAP 5th year data provided no evidence for cross-correlation in the large angular range ($l<60$). A signal at the 2--3 $\sigma$ level was however found at smaller scales, although its significance increased with increasing flux thresholds applied on NVSS sources (in contradiction with expectations for the ISW probed by NVSS and raising the issue of radio point source contamination). Furthermore, the intrinsic clustering of NVSS sources on the large scales (relevant for the ISW) was found too high for the commonly assumed redshift distribution for NVSS sources, as found in previous works, \citep{Negrello2006,Raccanelli2008}. Regarding ISW detection claims based upon SDSS data, there is also some ongoing discussion after recent failures in finding any statistical significance for the ISW, \citep{Bielby2008,Martin2010,Tom2010}. This situation is partially caused by the fact that the ISW is generated on the large angular scales and at moderate to high redshifts ($z\in [0.1,1.3]$). Deep galaxy surveys covering large fractions of the sky are hence required to sample the ISW properly, but those are not available yet (or not properly understood).\\ Ideally one would try to find the ISW contribution to CMB anisotropies by using CMB data {\em exclusively}. And in this context the thermal Sunyaev-Zel'dovich effect becomes of relevance. The thermal Sunyaev-Zel'dovich (tSZ) effect \citep{SZ72} results from the inverse Compton scattering of the CMB photons off the galaxy clusters electrons, and is expected to provide cosmological constraints on the normalisation at 8$h^{-1}$ Mpc of the density fluctuations power spectrum, $\sigma_8$, as well as on the amount of matter $\Omega_m$ and to a lower extent on the dark energy equation of state \citep[e.g.][]{BattyeWeller03}. Since both the tSZ and the ISW effects probe large scale structures and their evolution, a correlation is therefore expected between these two signals. {As \citet{chmSunyaev2005} pointed out, provided that the tSZ has a definite and well known frequency dependence, it is possible to combine different CMB maps obtained at different frequencies in search for a {\em frequency dependent} ISW--tSZ cross-correlation. The advantages of this approach are two folded: {\it (i)} only CMB data (obtained at different frequency channels) are required, and hence there is no need for using and characterising an external LSS catalog, and {\it (ii)} a better handle on systematics is provided since the ISW--tSZ cross correlation has a perfectly known frequency dependence that can be searched for in multiple channel combinations. This approach in experiments like {\it Planck} \footnote{{\it Planck} URL site:\\ {\tt http://www.sciops.esa.int/index.php?project=PLANCK}}, covering the whole sky in a wide frequency range, is well suited to separate the tSZ from other components present in the microwave range. After introducing in Section 2 the theoretical background, we study in Section 3 a combination of CMB maps at different frequencies that provides an unbiased estimate of the ISW-tSZ angular power spectrum. We also compute the expected significance for the ISW--tSZ cross correlation. In Section 4 we analyze how the large scale structure contributes to the tSZ and ISW auto power spectra in different redshift ranges. This allows us to define a strategy to optimize the S/N of the ISW--tSZ cross-correlation by applying a selective mask on galaxy clusters. We then investigate in Section 5 the limitations of our method due to the presence of galactic and extragalactic foregrounds and suggest some approaches to minimize their impact. We present our conclusions in Section 6.

The problem of component separation in multifrequency CMB observations for experiments like {\it Planck} has been a subject of active investigation \citep[e.g.,][]{stompor09,leach,eriksen08,juan07,vlad05}. The different contaminants (either galactic and extragalactic) show generally a different spectral dependence when compared to the CMB. On the low frequency side, foreground powerful in radio wavelenghts fall steeply and become the sub-dominant foregrounds at frequencies above $\nu \sim 70$ GHz, \citep{forewmap1}. Above this frequency, the presence of dust absorbing UV radiation and re-emitting it in the sub-millimeter and millimeter range constitutes the dominant contaminant. An experiment like {\it Planck}, with four high-angular resolution channels covering the 217 -- 857 GHz frequency range, should provide an accurate description of this foreground. With these data in hand, currently existing models describing the physics of dust emission should be improved further and accurate extrapolations to lower frequencies should be enabled. In our toy model describing the impact of the contaminants we assumed that the 353 GHz channel was a perfect (CMB -- free) dust tracer, and that the signal at low frequencies (100 or 143 GHz) was either due to dust or a white Gaussian noise signal. These assumptions may be overly optimistic (in regard to the 100 \& 217 GHz channels themselves), but however a generic combination of all nine channels (ranging from 30 GHz up to 857 GHz) are expected to provide an estimation of each of the components (CMB + foregrounds) whose residuals are expected to lie close to the level of $\sim 0.1$ ($\mu$K)$^2$ \citep{leach}. This is already the accuracy ballpark that our simple analysis proved to define the regime of detectability of the tSZ -- ISW cross correlation, and there may still be room for a more optimized channel combination oriented to unveil the particular tSZ--ISW cross correlation. The use of high galactic latitude HI maps as tracers of galactic cirrus could allow to lower the impact of those residuals at the level of a percent in $C_\ell$, (G. Lagache, private communication). \cite{NestorFernandez08} computed the cirrus power spectrum at different frequencies and different HI column densities. They found, in the case of fields that have a very low level of dust contamination, that the cirrus power spectrum at 217 GHz is of the order of 5 ($\mu$K$_{\rm RJ})^2$ at $\ell=10$ in units of $l(l+1)C_l/(2\pi)$. A one percent residual of the cirrus emission is thus smaller than 0.05 ($\mu$K$_{\rm RJ})^2$, i.e., 0.4 ($\mu$K$_{\rm CMB})^2$. Current foreground residual estimates based upon the works of \cite{NestorFernandez08} and \cite{leach} suggest that at the frequencies of interest (100 -- 217 GHz) the contaminant residuals remain a factor of a few above our requirements. How much room there is for improvement below those limits is something yet to be estimated from real data. {\it Planck} data at high frequencies provide a more profound knowledge of the dust properties in both our galaxy and extragalactic sources, together with the mechanisms involving its emission in the sub-millimeter range. It is nevertheless important to bear in mind that, since the different frequencies sample the infrared galaxy populations {at different redshifts} and the galaxy linear bias evolves with redshift \citep[e.g.][]{Lagache07,VieroBlast09}, using high frequency maps so as to clean the 143 and 217 GHz could potentially degrade the residual level we obtained in Section 5 with our template fitting method. This issue is under current investigation within the {\it Planck} collaboration. Even in the worst scenario in which foreground residuals are too high and complicated to prevent the detection of the tSZ -- ISW cross correlation, the upper limits to be imposed on it are of cosmological relevance, since it would constrain cosmological parameters like $\sigma_8$, $\Omega_m$ or $\Omega_{\Lambda}$. We have shown that the tSZ -- ISW cross correlation constitutes a CMB contained test for Dark Energy. The peculiar frequency dependence of the tSZ effect and the availability of multi-frequency all sky CMB observations provided by the experiment {\it Planck} should enable an estimation of this cross correlation provided the hot gas is a fair tracer of the potential wells during the cosmological epochs where the ISW is active. Our theoretical study shows that the Poisson/shot noise introduced by the modest number of very massive, very bright in tSZ galaxy clusters can be attenuated by masking out those tSZ sources below redshift $z< 0.3$. In the absence of a massive galaxy cluster catalog below that redshift, it would suffice to excise from the analysis those tSZ clusters clearly detected by {\it Planck} in order to achieve S/N of the order of 3.9 ($f_{sky}=1$). This tSZ -- ISW cross correlation detection would not require the use of deep-in-redshift and wide-in-angle galaxy surveys, but only the combination of different frequency CMB observations. This would hence provide a different approach for ISW detection with different systematics to other attempts based upon CMB -- galaxy survey cross correlations. If foreground residuals are kept at or below the $\sim 0.04 $ ($\mu$K)$^2$ level (in $l(l+1)C_l/(2\pi)$ units at $l\sim 10$) in the frequency range 100 -- 217 GHz, then the tSZ -- ISW correlation should provide a valid and independent test for the impact of Dark Energy on the growth of structure and the evolution of large angle CMB temperature anisotropies.

1012.5036

1012

1012.2420_arXiv.txt

We illustrate the structure and the main phenomenological features of a supersymmetric model (the USSM-A) built following a bottom-up approach and containing an anomalous abelian gauge symmetry. This model supports a gauged axion in its spectrum and provides a generalization of the global (supersymmetric) Peccei-Quinn construction. Complete simulations of the neutralino relic density are performed. Bounds from CAST and WMAP, combined with dark matter simulations, provide significant constraints on the scale of the interactions between the axion and the gauge fields. \vspace{1pc}

Gauged shift symmetries are typical of potentials containing flat directions and are common to several theories built around the Planck scale, such as strings/branes, down to supergravity and supersymmetric theories. The latter inherit the rich structure of brane models via some mechanism of dimensional reduction or geometric compactification, in the presence of external fluxes. However, the dynamics of scalar fields (massive and massless), associated to these flat directions, which take the form of massive moduli and massless Goldstone modes, is quite involved. This is in part due to the excessive proliferation of scalars in the low energy theory. On the other hand, one could also take into consideration the possibility that some of these flat directions could be (almost) preserved as the Universe expands, down to the low energy scale. In particular, they could be slightly lifted (non-perturbatively) from their unperturbed vacuum value at the electroweak and QCD phase transitions. In this case one could envision possible contributions to the dark matter density from these very light scalars/pseudoscalars. Although rather exotic at a first glance, this picture is not new, since it has been at the core of the long known Peccei-Quinn (PQ) proposal for the solution of the strong CP-problem, that shares some of its typical features. According to this mechanism, which invokes a very light pseudoscalar in the physical spectrum, a global anomalous $U(1)$ symmetry is attached to the fields of the Standard Model, which breaks at a very large scale ($f_a\sim 10^{12}$ GeV). In turn, the nature of the axion, as a (pseudo) Nambu-Goldstone mode of the broken global $U(1)$, has to be found at a later stage in the Early Universe, at the QCD transition ($\sim \Lambda_{QCD}$), with the generation of a periodic potential due to the instanton vacuum. The rather singular nature of this mechanism, which involves two widely separated scales, is made evident by the expression of the axion mass ($m_a$) which is related to their ratio $(m_a\sim\Lambda_{QCD}^2/f_a)$ and by its rather small value ($m_a\sim 10^{-3}$ eV). In a supersymmetric context, the axion is the imaginary component of a complex scalar ($b$), and is accompanied by another degree of freedom, the saxion, described by $\textrm{Re}\, b$, and by a supersymmetric fermionic partner (the axino, $\psi_b$). The gauging of this multiplet has been discussed in \cite{Kors:2004ri}, and further in \cite{Feldman:2010wy} in a study of the St\"uckelberg extension of the MSSM, with a single extra (non anomalous) $U(1)$ symmetry. In this extension $\textrm{Im}\, b$ is a Goldstone mode of the $U(1)$ gauge symmetry and, as such, becomes the longitudinal component of the $U(1)$ gauge field, disappearing completely from the physical spectrum. A similar destiny is shared by the axion of a $U(1)'$ (prime) extension of the MSSM (U(1)' MSSM) \cite{Anastasopoulos:2008jt}, which is anomalous and introduces extra interactions in the form of PQ counterterms for the restoration of gauge invariance (supersymmetric terms which generalize the $\textrm{Im}\, b\, F\tilde{F}$ vertex). In this second model the axino mixes with the gauginos and Higgsinos of the theory to generate several neutralinos and the corresponding LSP (light supersymmetric particle). The first model that supports a physical axion and is compatible with supersymmetry, called the USSM-A, is built around the structure of the USSM \cite{Coriano:2008aw,Coriano:2008xa}, with some important variants: 1) there is no extra Higgs to ensure the breaking of the $U(1)$ symmetry, rather, this is realized in the St\"uckelberg form; 2) the $U(1)$ symmetry is anomalous. These two conditions, together with a requirement on the charge assignments of the Higgs sector that mixes the St\"uckelberg and the Higgs mechanisms, allow to construct a complete supersymmetric model where the gauged axion is a component of $\textrm{Im}\, b$. \subsection{The physical axion} The extraction of a physical axion in theories of this type was pointed out in \cite{Coriano:2005js} in a non-supersymmetric context, motivated in a certain class of string vacua, and in successive phenomenological studies \cite{Coriano:2007fw,Coriano:2007xg}. It was later remarked that these effective actions could be generated starting from anomaly-free theories, with no reference to any class of string vacua, under particular conditions on the decoupling of a chiral fermion. An example of such a behaviour is encountered when a heavy Higgs decouples from the low energy spectrum, "dragging" away also a part of the fermion spectrum (via Yukawa interactions) and one gauge boson, which becomes very massive \cite{Coriano:2009zh}. The effective low energy theory carries the signature of this "partial decoupling" of the Higgs, inheriting effective interactions which are PQ like. In this case, as in the original PQ model, the low energy axion is related to the phase of the heavy Higgs field. Before coming to the issue of the mass of the axion in this model, we should mention that this is generated by an extra potential, allowed by the gauge symmetry \cite{Coriano:2005js}, which is periodic in the physical axion field \cite{Coriano:2010ws,Coriano:2010py}. The potential is gauge invariant only if the the extra singlet superfield $\hat{S}$ of the superpotential is charged under the anomalous $U(1)$. In turns, this implies that the two Higgs superfields of the model are also charged under the same $U(1)$ and are not charge-aligned. As a result of this choice, the mass of the anomalous gauge boson is induced both by the Higgs and the St\"uckelberg mechanisms. We briefly comment on other essential features. 1) The anomalous fermion spectrum induces trilinear gauge interactions which are absent in anomaly-free extensions. These consist of cubic $(U(1)_B^3)$ anomalies and mixed-anomalies of $U(1)_B$ with the hypercharge gauge field (Y), and in combination with non-abelian anomalies with the $SU(2)$ and $SU(3)$ gauge fields. 2) For a general anomalous interaction, the D-terms of the scalar potential are non-local, but expandable in the St\"uckelberg scale. This feature can be avoided by a suitable choice of the fermion charges, without causing a complete decoupling of the extra $U(1)_B$. 3) The (St\"uckelberg) mass of the anomalous gauge boson, $M_{St}$, is a free parameter of the theory and is the suppression scale of the $\textrm{ Im}\, b\, F \tilde F$ interactions.

1012.2420

1012

1012.0755_arXiv.txt

We have performed the first measurement of the angular power spectrum in the large­-scale diffuse emission at energies from 1-­50 GeV. We compared results from data and a simulated model in order to identify significant differences in anisotropy properties. We found angular power above the photon noise level in the data at multipoles greater than $\sim 100$ for energies $1 \lesssim E \lesssim 10$ GeV. The excess power in the data suggests a contribution from a point source population not present in the model.

The Fermi Gamma-Ray Telescope, launched on June 11th 2008 from Cape Canaveral, performs gamma-ray measurements over the whole celestial sphere. Its main scientific instrument the Large Area Telescope (LAT) measures the tracks of the electron and positron that result when an incident gamma-ray undergoes pair conversion in a thin, tungsten foil, and measures the energy of the subsequent electromagnetic shower that develops in the telescope's calorimeter. Some \textit{Fermi}-LAT specifications are: Energy range from 20MeV to $\sim$300GeV, angular resolution $\sim$0.1 deg above 10 GeV, field of view (FoV) $\sim$2.4sr, and uniform sky exposure of $\sim$30 minutes every 3 hours. Detailed descriptions of the \textit{Fermi}-LAT telescope can be found in \cite{Atwood:2009ez}.% One of the key science targets of the Fermi mission is diffuse gamma-ray emission. Its main component is correlated with Milky Way structures, the galactic emission, arising from interactions of high-energy cosmic rays with the interstellar medium and the interstellar radiation field. A fainter component considered to have an isotropic or nearly isotropic distribution on the sky, the so-called extragalactic emission, has been observed. This observation is based on the modelization of galactic component, detected \textit{Fermi}-LAT sources and the solar gamma-ray emission \cite{Abdo:2010nz}. Also there is a contribution from populations of sources, of various kinds, including blazars, pulsars, SNR, and possibly dark matter (DM) structures, not yet detected due to \textit{Fermi}-LAT spatial resolution and photon statistics. The angular distribution of photons in the diffuse gamma-­ray background contains information about the presence and nature of these unresolved source populations (USP). Fluctuations on small scales may originate from USP if they are different from those expected from the Poisson noise due to finite statistics.% Recent studies have predicted the contributions to the angular power spectrum (APS) from extragalactic and galactic DM annihilation or decay, e.g. \cite{Ando:2005xg, Ando:2006cr, SiegalGaskins:2008ge, SiegalGaskins:2009ux, Ando:2009fp,Fornasa:2009qh,Zavala:2010pw,Taoso:2008qz}. A detailed \textit{Fermi}-LAT sensitivity study of anisotropies from DM annihilation has been presented in \cite{Cuoco:2010jb}.% I present the results of an anisotropy analysis of the diffuse emission measured by the \textit{Fermi}-LAT. We calculate the angular power spectrum of the emission from $\sim 22$ months of Fermi data and of the emission from a simulated model (galactic diffuse emission, 11-month sources from Fermi catalog and isotropic emission), and compare the results from the data and model in order to identify significant differences in anisotropy properties.

Plots of Fig. 1. show the APS of the data and the default model (Galactic diffuse model + 11 month source catalog + isotropic) in different energy ranges. These figures show at what energy ranges and multipole ranges the APS of the data and the model differ, as well as where each of these is consistent with the photon noise level.% We have found that at multipoles greater than $\sim 100$ the excess power in the data suggest a contribution from a point source population not present in the model. Also, at large angular scales ($l<100$) angular power above the noise is seen in the data and model, probably due to contamination from the galactic diffuse.% Due to decreasing photon statistics, the amplitude of anisotropies detectable by this analysis decreases with increasing energy. For this reason, the non-detection of power above the noise level at $10$-$­50$ GeV does not exclude the presence of anisotropies at the level of those detected at $1$-­$10$ GeV. \begin{center} \begin{figure}[htb]\label{plots} \includegraphics[width=0.45\textwidth]{german_gomez_fig01}\quad \includegraphics[width=0.45\textwidth]{german_gomez_fig02}\quad \includegraphics[width=0.45\textwidth]{german_gomez_fig03}\quad \includegraphics[width=0.45\textwidth]{german_gomez_fig04}\quad \includegraphics[width=0.45\textwidth]{german_gomez_fig05} \caption{Fluctuation APS of the data and the default model (Galactic diffuse model + 11 month source catalog + isotropic) in different energy ranges. } \end{figure} \end{center}

1012.0755

1012

1012.3083_arXiv.txt

In this contribution we summarize two recent applications of a correspondence between backreaction terms in averaged inhomogeneous cosmologies and an effective scalar field (the ``morphon''). Backreaction terms that add to the standard sources of Friedmannian kinematical laws and that emerge from geometrical curvature invariants built from inhomogeneities, can be interpreted in terms of a minimally coupled scalar field in the case of a dust matter source. We consider closure conditions of the averaged equations that lead to different evolution scenarii: a) the standard Chaplygin equation of state imposed as an effective relation between kinematical fluctuations and intrinsic curvature of space sections, and b) an inflationary scenario that emerges out of inhomogeneities of the Einstein vacuum, where averaged curvature inhomogeneities model the potential of an effective classical inflaton.

Our universe is supposed to verify the strong cosmological principle which demands homogeneity and isotropy at all scales. This standard approach, known as Friedmann--Lema\^{\i}tre--Robertson--Walker (FLRW) cosmology obeys the set of equations, \begin{subequations} \begin{eqnarray} \left(\frac{\dot a}{a}\right)^2 = \, \frac{8 \pi G}{3} \sum_i \varrho_{(i)} \, - \, \frac{k}{a^2} \, , \\ \frac{\ddot a}{a} \, = \, - \frac{4 \pi G}{3} \sum_i \left(\varrho_{(i)} + 3 p_{(i)}\right) \, , \\ \dot{\varrho}_{(i)} \, + \, 3 \, \frac{\dot a}{a} \left( \varrho_{(i)} + p_{(i)}\right) \, = \, 0 \, , \label{eq:hom_dyn} \end{eqnarray} \end{subequations} that, together with equations of state as relations between the homogeneous variables, e.g. between the pressures and energy densities, forms a closed system. The FLRW framework is widely used in order to describe the dynamics of our Universe and the formation of its constituents. It, however, leaves in suspense an explanation about the origin of Dark Energy and Dark Matter, which respectively represent in this model about 3/4 and 1/4 of the total content of the universe model. This last point might actually reveal a symptom of a deeper problem linked to this approach. Indeed, in FLRW cosmology one determines background quantities regardless of the scale and makes them evolve according to a homogeneous--isotropic solution of Einstein's equations. But, are the background quantities well--defined within standard cosmology, {\it i.e.} as a suitable average over the inhomogeneities? Is their evolution well--approximated in this framework, {\it i.e.} is the time dependence of the homogeneous--isotropic averaged state well approximated by a homogeneous--isotropic solution? All these issues can be routed back to the nonlinearity of the gravitational equations and the averaging problem \cite{ellis:average,ellisbuchert,buchert:review}. Rewriting Einstein's equations via the ADM formalism and spatially averaging, in a domain--dependent way, the scalar parts of the equations with respect to free--falling fluid elements, the averaged dynamics of an inhomogeneous universe model filled with a pressureless fluid assumes the following form \cite{buchert:dust,buchert:fluid,buchert:review}: \begin{subequations} \begin{eqnarray} \left(\frac{{\dot a}_\CD}{a_\CD}\right)^2 = \, \frac{8 \pi G}{3} \average{\varrho} \, - \, \frac{k_\initial\CD}{a_\CD^2} - \frac{\CW_\CD + {\CQ}_\CD}{6} \, , \\ \frac{{\ddot a}_\CD}{a_\CD} \, = \, - \frac{4 \pi G}{3} \average{\varrho} \, + \, \frac{{\CQ}_\CD}{3} \, , \\ \langle{\varrho}\rangle\dot{}_\CD \,+ \, 3 \, \frac{{\dot a}_\CD}{a_\CD} \average{\varrho} \, = \, 0 \, , \\ \frac{1}{a_\CD^6} \, \left(\,{\CQ}_\CD \, a_\CD^6 \,\right)\dot{} \, + \, \frac{1}{a_\CD^{2}} \, \left(\CW_\CD \, a_\CD^2 \, \right)\dot{} \, = \, 0 \, , \label{eq:inhom_dyn} \end{eqnarray} \end{subequations} where $a_\CD$ is the effective scale factor, $\average{\varrho}$ is the energy density of the irrotational dust averaged over a compact, mass--preserving domain $\CD$, $\CQ_\CD$ is the kinematical backreaction term (an extrinsic curvature invariant), and finally $\CW_\CD$ is the curvature deviation from a constant--curvature $k_{\initial\CD}$ according to the FLRW solution (an intrinsic curvature invariant). These variables are defined as \begin{subequations} \begin{eqnarray} \average{\varrho}(t) = \frac{M_{\initial\CD}}{V_{\initial\CD}} \, a_\CD^{-3} \,\, , \,\, a_\CD (t) := \left( \frac{V_{\CD}(t)}{V_{\initial\CD}} \right)^{1/3} \,, \\ {\CQ}_\CD (t) := \frac{2}{3}\average{\left(\theta - \average{\theta}\right)^2 } - 2\average{\sigma^2} \, , \\ \CW_\CD (t) := \average{\CR} - \frac{6 k_\initial\CD}{a_\CD^2}\, , \end{eqnarray} \end{subequations} with $V_{\initial\CD}$ the initial volume of the domain and $V_{\CD}(t)$ its volume at a proper time $t$, $\theta$ the rate of expansion, $\sigma$ the rate of shear, and $R$ the $3-$Ricci scalar curvature. Comparing the set of equations (1) and (2), one easily notices that the average evolution of an inhomogeneous universe model differs from the evolution of a homogeneous one. The change of the cosmological background evolution is here driven by the non--trivial geometrical structure of an inhomogeneous space, and the corresponding deviations are encoded into the backreaction terms $\CQ_\CD$ and $\CW_\CD$, which are coupled through the relation (\ref{eq:inhom_dyn}). Other formulations of the effective inhomogeneous equations can be done. The first one aims at considering the backreaction terms as an effective fluid by introducing an effective perfect fluid energy momentum tensor with \begin{equation} \varrho^{\CD}_{b} = - \frac{1}{16\pi G} ({\CQ}_\CD + \CW_\CD )\;;\; p^{\CD}_{b} = - \frac{1}{16\pi G} ({\CQ}_\CD - \frac{1}{3} \CW_\CD ), \label{backreactionfluid} \end{equation} leading to the following reformulation of the system (2): \begin{subequations} \begin{eqnarray} \left(\frac{{\dot a}_\CD}{a_\CD}\right)^2 \, = \, \frac{8 \pi G}{3} \, \left( \average{\varrho} + \varrho^{\CD}_{b} \right) \, - \, \frac{k_{\CD_{\it i}}}{a^2_\CD} \, , \\ \frac{{\ddot a}_\CD}{a_\CD} = \, - \frac{4 \pi G}{3} \, \left( \average{\varrho} + \varrho^{\CD}_{b} + 3 {p}^{\CD}_{b} \right) \, , \\ \langle{\varrho}\rangle\dot{}_\CD \,+ \, 3 \, \frac{{\dot a}_\CD}{a_\CD} \average{\varrho} \, = \, 0 \, , \\ {\dot\varrho}^{\CD}_{b} + 3 \, \frac{{\dot a}_\CD}{a_\CD} \, \left(\varrho^{\CD}_{b} + {p}^{\CD}_{b} \right) = 0 \, , \end{eqnarray} \end{subequations} where we see that the coupling between the backreaction terms now stands for the conservation law of the backreaction fluid. We also appreciate that, like in the standard model, we need an equation of state to close the system that here is dynamical and furnishes a relation between the effective sources. We shall consider two of such equations of state employed as closure conditions for the set of equations (5). The second reformulation, suggested by the form of the effective sources (\ref{backreactionfluid}) \cite{buchert:static}, consists in describing the backreaction fluid by a minimally coupled real scalar field $\phi_\CD$,~called the {\it morphon field}, evolving in an effective potential $U_\CD(\phi_\CD)$ \cite{morphon}: \begin{equation} \varrho_{b}^{\CD} = \varrho_{\phi}^{\CD} : = \epsilon \dot{\phi}^2_{\CD}/2 + U_{\CD}\;;\;p_{b}^{\CD} = p_{\phi}^{\CD} : = \epsilon \dot{\phi}^2_{\CD}/2 - U_{\CD}, \end{equation} where $\epsilon = + 1$ for a standard scalar field (with a positive kinetic energy), and $\epsilon = - 1$ for a phantom scalar field (with a negative kinetic energy). The scalar field language leads to the following {\it correspondence relations}: \begin{subequations} \begin{eqnarray} \epsilon \dot{\phi}^2_{\CD} = - \frac{1}{8 \pi G} \, ({\CQ}_\CD + \frac{1}{3}{\CW}_\CD ) \, , \\ U_{\CD} = - \frac{1}{24 \pi G} {\CW}_\CD \, , \\ \ddot{\phi}_{\CD} + 3 \frac{\dot{a}_\CD}{a_\CD} \, \dot{\phi}_{\CD} + \epsilon \frac{\partial}{\partial \phi_{\CD}} U_{\CD} = 0 \, , \label{eq:morphon} \end{eqnarray} \end{subequations} where the Klein--Gordon equation is simply the counterpart of the conservation law for the backreaction fluid. This correspondence allows us to interpret the kinematical backreaction effects in terms of the properties of scalar field cosmologies, notably quintessence or phantom--quintessence scenarii that are here routed back to models of inhomogeneities.

1012.3083

1012

1012.4801_arXiv.txt

Using \suzaku and the \textsl{Rossi X-ray Timing Explorer} (\rxte), we have conducted a series of four simultaneous observations of the galactic black hole candidate \cyg in what were historically faint and spectrally hard ``low states''. Additionally, all of these observations occurred near superior conjunction with our line of sight to the X-ray source passing through the dense phases of the ``focused wind'' from the mass donating secondary. One of our observations was also simultaneous with observations by the \chandra-High Energy Transmission Grating (\hetg). These latter spectra are crucial for {revealing} the ionized absorption due to the secondary's focused wind. Such absorption is present and must be accounted for in all four spectra. These simultaneous data give an unprecedented view of the 0.8--300\,keV spectrum of \cyg, and hence bear upon both corona and X-ray emitting jet models of black hole hard states. Three models fit the spectra well: coronae with thermal or mixed thermal/non-thermal electron populations, and jets. All three models require a soft component that we fit with a low temperature disk spectrum with an inner radius of only a few tens of $GM/c^2$. All three models also agree that the known spectral break at 10\,keV is not solely due to the presence of reflection, but each gives a different underlying explanation for the augmentation of this break. Thus whereas all three models require that there is a relativistically broadened Fe line, the strength and inner radius of such a line is dependent upon the specific model, {thus making premature line-based estimates of the black hole spin in the Cyg X-1 system}. We look at the relativistic line in detail, accounting for the narrow Fe emission and ionized absorption detected by \hetg. Although the specific relativistic parameters of the line are continuum-dependent, none of the broad line fits allow for an inner disk radius that is $>40$\,$GM/c^2$.

\label{sec:intro} \setcounter{footnote}{0} There is currently significant debate as to the physical mechanisms responsible for the continuum of X-ray spectrally ``hard states'' of black hole candidates (BHC) accreting in binaries. This debate ranges from the broader issue of whether or not there is a significant contribution to the X-ray band from an outflow or jet, to more narrowly focused issues within given classes of models. For instance, the hard X-ray emission has traditionally been attributed to a Comptonizing thermal corona \citep[][and references therein]{eardley:75a,shapiro:76a,sunyaev:79a,dove:98a}. These earlier works generally favor a scenario where the corona lies central to a truncated outer thin disk. However, if the corona is driven outwards by radiative pressure \citep[e.g.,][] {beloborodov:99a} could it instead overlay the inner disk? Can the optically thick, geometrically thin disk extend inward nearly to the innermost stable circular orbit \citep{miller:06b}? Is this disk cold (peak temperatures of a few hundred eV), or can it instead be hot (near a keV, i.e., \citealt{wilms:06a})? Is the hard state corona comprised primarily of electrons with a thermal population \citep{poutanen:09a}, or can it have a substantial contribution from a non-thermal electron population \citep{ibragimov:05a}? Does the bulk motion of the flow play a role in Comptonizing the spectrum \citep{shaposhnikov:06a,laurent:07a}? Alternatively could the X-rays be comprised of a combination of direct synchrotron and synchrotron self-Compton (SSC) emission from a jet \citep{markoff:05a,maitra:09a}? Contributing to the debate, however, is the fact that many of the above cited models, especially when considering solely \textsl{Rossi X-ray Timing Explorer} (\rxte) data in the 3--200\,keV range (or even narrower energy ranges), describe the hard state spectra nearly equally well. To study these issues, over the past decade we have been using a series of pointed, approximately bi-weekly, \rxte observations of the BHC \cyg performed simultaneously with 15\,GHz radio observations by the Ryle telescope. \cyg holds much promise for exploring the current range of questions listed above owing to its persistently bright X-ray flux (including both the hard and soft states, it varies between 200--600\,mCrab in the 1.2--12 keV band covered by the \rxte-All Sky Monitor) and its correlated radio/X-ray spectra \citep[see][and references therein]{wilms:06a}. Extended radio emission even has been imaged in Cyg X-1 \citep{stirling:01a}. This campaign has already provided the spectra for some of the Comptonization models \citep{wilms:06a}, and jet models \citep{markoff:05a} discussed above. Additionally, these data have been used to to study the correlation of the X-ray and radio spectral properties on both long time scales \citep{pottschmidt:00a,pottschmidt:02a,wilms:06a}, and short time scales (\citealt{gleissner:04a}, \citealt{wilms:07a}, \citealt{boeck:10a}, B\"ock et al., in prep.). Thus they comprise a strong data set for addressing not only the details of the X-ray spectrum, but also the connection to the hard state jets which are known to dominate the radio through near-infrared emission. Simultaneous radio/X-ray flaring has also been detected in \cyg during this extended campaign (\citealt{fender:06a}, \citealt{wilms:07a}), lending support to the hypothesis of X-ray emission by the jet. \rxte spectral data alone, however, do not allow us to break the current existing ``theoretical degeneracy'' in the origin of the X-ray spectrum. The statistically best fits are in fact obtained with purely empirical simple broken powerlaws with a break occurring between 9--12\,keV, and an exponential cutoff occurring at >20\,keV, to which a broad, $\approx 6.4$\,keV gaussian line is added \citep{wilms:06a,nowak:05a}. The latter component is likely attributable to a relativistically broadened Fe K$\alpha$ line \citep[][and references therein]{reynolds:03a}; however, its parameters are dependent upon the assumed continuum model \citep{wilms:06a}. More physically motivated Comptonization and outflow-dominated models \citep{markoff:05a,wilms:06a} can describe the same spectra almost as well. However they must introduce additional, albeit plausible, physical components (e.g., relativistic smearing) to recover the simple spectra of the broken powerlaw description. Thus, there is some amount of ambiguity when correlating detailed spectral features vs. continuum properties, e.g., disk reflection vs. coronal compactness/hardness, as the detailed features systematically depend upon the underlying broad-band continuum. Despite these problems in finding a truly unique spectral model, when considering multiple observations taken over a wide range of luminosities and spectral hardnesses, spectral correlations arise that are robust and persistent across a variety of these theoretical characterizations \citep{wilms:06a}. Using the broken powerlaw models as a simple description of the X-ray spectra, we have found that when the 2--10\,keV photon index $\Gamma_1 < 2.2$ there is a positive correlation between X-ray and radio flux, whereas for $\Gamma_1 > 2.2$ there is a negative correlation between the X-ray and radio flux. We therefore use the value of $\Gamma_1 \approx 2.2$ as the canonical division between the spectrally ``hard state'' and the spectrally ``soft state'' \citep[see][]{remillard:06a}. Additionally, as the spectra become harder, exponential cutoffs tend to become less significant \citep[see also][]{motta:09a}. Our previous \rxte spectral studies of \cyg have been limited in two respects: the low spectral resolution of \rxte ($E/\Delta E \approx 6$ at 6\,keV), and the inability to measure spectra at $\aproxlt 3$\,keV. In this work, we turn to a set of four \suzaku observations that we performed simultaneously with our \rxte campaign to enhance our \cyg studies in several crucial ways. First, \suzaku has large effective area at soft X-ray energies. In this work, we consider spectra down to 0.8\,keV, which allows for the possibility of measuring the ``seed photon spectrum'' in Comptonization models (\S\ref{sec:compton}), or judging the relative contribution of synchrotron versus the disk radiation in jet models (\S\ref{sec:jet}). \suzaku also has excellent resolution in the Fe K$\alpha$ line region ($E/\Delta E \approx 50$), which allows us to separate narrow from relativistically broadened line features (\S\ref{sec:comp_line}, \ref{sec:rel_line}). Third, \suzaku measures the \cyg hard X-ray spectrum up to $\approx 300$\,keV (\S\ref{sec:cutoff}), providing further constraints on Comptonization and jet models. \begin{deluxetable}{ccccc} \setlength{\tabcolsep}{0.03in} \tabletypesize{\footnotesize} \tablewidth{0pt} \tablecaption{Log of \cyg Observations \label{tab:lo}} \tablehead{\colhead{Date} & \colhead{Spacecraft/ObsID} & \colhead{Instrument} & \multicolumn{2}{c}{Exposure} \\ (yyyy-mm-dd) & & & \multicolumn{2}{c}{(ksec)} } \startdata 2006-10-30 & \suzaku/401059010 & \xis0--3 & 35.0\tablenotemark{a} & 32.1 \\ \nodata & \nodata & \hxd-\pin & 27.7 & 24.9 \\ \nodata & \nodata & \hxd-\gso & 27.7 & 25.8 \\ \nodata & \rxte/80110-01-13 & \pca & 8.0\tablenotemark{b} & 8.0 \\ \nodata & \nodata & \hexte-A & 3.1 & 3.1 \\ \nodata & \nodata & \hexte-B & 2.5 & 2.5 \\ 2007-04-30 & \suzaku/402072010 & \xis0,1,3 & 34.0\tablenotemark{a} & 22.4 \\ \nodata & \nodata & \hxd-\pin & 40.2 & 27.6 \\ \nodata & \nodata & \hxd-\gso & 40.2 & 27.1 \\ \nodata & \rxte/92090-01-16 & \pca & 14.6\tablenotemark{b} & 1.5 \\ \nodata & \nodata & \hexte-A & 4.6\tablenotemark{c,d} & 0.3 \\ \nodata & \nodata & \hexte-B & 3.0 & 0.4 \\ 2007-05-17 & \suzaku/402072020 & \xis0,1,3 & 22.3\tablenotemark{a} & 12.3 \\ \nodata & \nodata & \hxd-\pin & 32.6 & 10.9 \\ \nodata & \nodata & \hxd-\gso & 32.6 & 10.9 \\ \nodata & \rxte/92090-01-17 & \pca & 2.0\tablenotemark{b} & 0.0 \\ \nodata & \nodata & \hexte-A & 1.9\tablenotemark{c,d} & 0.0 \\ \nodata & \nodata & \hexte-B & 0.9 & 0.0 \\ 2008-04-19 & \suzaku/403065010 & \xis1,3 & 16.9\tablenotemark{a} & 7.2 \\ \nodata & \nodata & \xis0 & 34.0\tablenotemark{e} & 0.0 \\ \nodata & \nodata & \hxd-\pin & 29.0 & 10.1 \\ \nodata & \nodata & \hxd-\gso & 29.0 & 10.1 \\ \nodata & \rxte/93120-01-01 & \pca & 21.5\tablenotemark{b} & 7.2 \\ \nodata & \nodata & \hexte-A & 17.9\tablenotemark{c} & 7.1 \\ \nodata & \nodata & \hexte-B & 10.8 & 4.6 \\ \nodata & \chandra/8525 & \hetg & 29.4\tablenotemark{f} & 11.1\tablenotemark{g} \\ \enddata \tablecomments{Exposure times are after initial good time filtering {(left), and after color/intensity time filtering (right)}.} \tablenotetext{a}{Summed exposure times for all listed \xis detectors.} \tablenotetext{b}{Summed exposure intervals, not weighted by the fraction of operating Proportional Counter Units (PCU).} \tablenotetext{c}{HEXTE-A cluster in fixed position (no rocking).} \tablenotetext{d}{Evidence for HEXTE-A cluster exposure time anomaly.} \tablenotetext{e}{\xis0 run in continuous readout mode.} \tablenotetext{f}{Summed exposure times for all $1^{\rm st}$ order spectra.} \tablenotetext{g}{Summed exposure times for $1^{\rm st}$ order \heg spectra.} \end{deluxetable} The outline of this paper is as follows. In \S\ref{sec:data}, we describe our full set of observations and our data reduction procedures. Due to maintenance and upgrade of the Ryle radio telescope during the construction of the Arcminute Microkelvin Imager (AMI), no simultaneous radio measurements are available for these \suzaku/\rxte observations. Instead, we use previous observations to estimate the radio fluxes (\S\ref{sec:radio}). For one of our observations simultaneous \chandra-\hetg data are available (\S\ref{sec:data_chandra}). These \chandra data become crucial in all of our analyses as they help elucidate the spectral variability associated with the observed lightcurve behavior, as discussed in \S\ref{sec:lightcurves}. We present simple phenomenological descriptions of the spectra in \S\ref{sec:simple}, including a description of the composite line profile (\S\ref{sec:comp_line}). Comptonization and jet models are presented in \S\ref{sec:bb}, along with further discussions of the implied relativistic lines (\S\ref{sec:rel_line}). We summarize our findings in \S\ref{sec:discuss}.

\label{sec:discuss} In this work we present broad band (0.8--300\,keV) fits to four separate observations of \cyg that have simultaneous \suzaku and \rxte data. The most recent of these observations also has simultaneous \chandra-\hetg data. Each of these observations shows evidence of dipping events likely due to dense structures (``clumps'') in the otherwise highly ionized wind of the secondary. This is seen in the lightcurves (Figs.~\ref{fig:lc}), and the color-color diagrams (Fig.~\ref{fig:colors}, which {can be modeled by} the presence of a dust scattering halo). The ionized absorption is very significant in the \suzaku spectra (Fig.~\ref{fig:suzaku_chandra}) even though in this work we consider the least absorbed {periods} of the lightcurve. \begin{itemize} \item Modeling the ionized line absorption present in \cyg, via the use of parameterized fits to the \hetg spectra, is crucial for deriving good fits to the soft X-ray spectra seen by \suzaku and \rxte. \end{itemize} We note that even accounting for this ionized absorption, the spectral fits presented in this paper yield reduced $\chi^2_\nu$ that range from 1.8--2.5. Do such values truly represent good fits to these data? \cyg is bright, and these observations are of sufficient length that the signal-to-noise values for these spectra are quite high. The spectra are dominated by systematic errors, especially at the soft X-ray energies. We already have added 0.5\% systematic errors to the \pca spectra, which is a reasonable estimate for the \emph{internal} uncertainty of the \pca. There are also relative uncertainties among the detectors. For observation 4, if we increase the \pca systematic errors to 1\%, add 1\% systematic errors to the \hexte A cluster (i.e., the fixed cluster), and add 3\% systematic errors to the \suzaku-\xis spectra, then the fits presented here would have reduced $\chi^2_\nu \approx 1$. Comparing the fits among the individual \xis spectra (representing both different individual detectors and different data acquisition modes), we have found that $\pm 3\%$ is a reasonable estimate of the end-to-end differences among these spectra. We hypothesize that the quality of the fits presented here is near ``optimal'' given the current internal and relative calibrations of these detectors. {The observations that occurred on 2007 May 17, i.e., the third set of observations, are potentially the most problematic in turns of cross-calibration issues. For our other sets of observations, the \rxte-\pca data act as a ``bridge'' between the \suzaku-\xis and -\pin spectra. Each can have its relative normalization anchored by a comparison to the \pca, which overlaps the energy coverage of both detectors. This is lacking in the third observation, leading to the worry that changes in the normalization constant are subsuming, for example, continuum spectra associated with the spectral break at $\approx 10$\,keV. For the third observation, we find the ratio of the \pin to \xis1 normalization constant to range from 1.10--1.17. This is to be compared to the 1.16--1.18 value found for fits to the first observation, the 1.06--1.08 found for the second observation, and the 1.10--1.11 found for the fourth observation. The expected value\footnote{{\tt http://heasarc.gsfc.nasa.gov/docs/suzaku/analysis/} {\tt watchout.html.} Since we did not fit \xis0 data for the fourth observation, here we compare to \xis1.} for the \pin/\xis0 comparison is 1.16. There is some amount of scatter for the cross-normalization values among the different observations; however, the first, second, and fourth observations show little scatter in its value for different fits to the same data. More scatter is seen for the third observation; therefore, additional systematic uncertainties need to be considered as being present for that set of observations.} Using simple broken powerlaw and exponentially cutoff powerlaw fits, we find that these spectra are among the hardest seen in the ``low hard state'' of \cyg over the past decade. For all four observations, the spectra are clearly detected out to 300\,keV with the \gso, and exponential cutoffs are well constrained. \begin{itemize} \item Although these are among the hardest \cyg spectra ever detected, the exponential folding energies vary by over a factor 1.5, and range from 160--250\,keV. \end{itemize} Historically observed hard state spectra in \cyg have shown folding energies that vary over a slightly wider range, while hard state BHC as a class show folding energies that span a factor of 5. As these spectra are among the faintest and hardest for \cyg, they make excellent test beds for theoretical models that posit, for instance, that the hard state represents a configuration with an inner disk that has evaporated into a quasi-spherical corona. In such a scenario, we might expect these spectra to show the most ``extreme'' recession of the inner disk, although we have noted that the bolometric luminosities represented by these spectra only span a factor of two. At a few percent of the Eddington luminosity, they are not far below the expected soft-to-hard state transition. Numerous transient BHC sources show much fainter hard states as they fade into quiescence. We have presented a number of different spectral models, all of which describe the 0.8--300\,keV spectra well. Some of these models describe the X-ray spectra primarily with Comptonization components (whether due to a thermal or hybrid thermal/non-thermal corona), while the jet model is dominated by synchrotron and SSC emission from the jet. All of these models have a number of features in common. \begin{itemize} \item All models require a soft excess that here we describe with a disk component with low ($kT_{\rm disk} \approx 200$\,eV) peak temperature. The implied inner radii of these disks range from 2--40\,$GM/c^2$. \item All models require a relativistically broadened line component. The inner emission radius of this broadened line never exceeds $\approx 40~GM/c^2$, but for some models is as low as 6\,$GM/c^2$. \item All models require a reflection component. The typical values for the reflection fraction are $\Omega/2\pi \approx 0.2$--0.3. \item All models imply that the spectral hardening at $\approx 10$\,keV is not \emph{solely} due to reflection. \end{itemize} This latter point is very important, and broadly agrees with similar conclusions drawn by \citet{frontera:01a}, \citet{ibragimov:05a}, and \citet{makishima:08a}. The presence of the broad Fe line and \emph{some} of the spectral curvature in the 20--300\,keV band is a clear indication of the presence of reflection. However, this reflection spectrum is \emph{not} sitting on top of a simple ``\texttt{disk+powerlaw}'' spectrum. There is additional continuum complexity separate from reflection that contributes to this perceived break. A high seed photon temperature in the thermal corona model yields a soft excess in the 2-10\,keV band and thus contributes to the measured break in that scenario. (See the discussion in \citealt{wilms:06a}.) As has been discussed by \citet{ibragimov:05a}, a non-thermal electron population in the corona can lead to a soft excess in the 2--10\,keV band that helps contribute to what otherwise would be modeled as a reflection break at 10\,keV. The two corona model discussed by \citet{makishima:08a}, a version of which is shown in Fig.~\ref{fig:deqpair}, rather explicitly replaces part of the reflection component with a broad band continuum model. Finally, in the jet paradigm the spectral break at 10\,keV is partly attributable to the transition from dominance by synchrotron emission to SSC emission in the continuum. There are plausible physical scenarios for each of the discussed spectral models, {with some hypothesized geometries being shown in Fig.~\ref{fig:geometry}}. The hybrid thermal/non-thermal coronal model is the closest to the concept of the quasi-spherical inner corona with outer geometrically thin disk \citep[e.g.,][etc.]{eardley:75a,shapiro:76a,ichimaru:77a,dove:97b}. We have only been able to find such solutions, however, when invoking a partly non-thermal electron population in the corona. The purely thermal coronal model contains two disk components that are reminiscent of the physical description given by \citet{mayer:07a}. These authors describe a situation where an outer, geometrically thin, cool disk surrounds an inner, geometrically thick, hot corona. In the very inner radii of this corona, however, thermal conduction leads to it condensing into a geometrically thin and optically thick disk {(see Fig.~\ref{fig:geometry})}. Such a component could supply the high temperature seed photons in our thermal corona solutions, {while the lower temperature outer disk could provide the bulk of the reflection features}. The jet model has a natural physical interpretation in that the usually observed optically thick radio spectrum observed in the hard state is clear indication of the presence of a jet. The question that remains is the contribution of this component to the X-ray band. The fact that the continuum is more complex than a simple ``\texttt{disk+powerlaw}'', yet there are multiple, physically motivated models that yield comparably good spectral fits, leads to the final point. \begin{itemize} \item Although a relativistically broadened line is required in all of our spectral models, the parameters of this line are dependent upon the presumed continuum model. \end{itemize} {Coupled with this dependence upon assumed continuum spectrum is an implicit dependence upon ionized absorption, for which we have detailed \chandra-\hetg measurements for only the fourth observation\footnote{Since performing these observations, we have carried out a \chandra-\hetg observation of orbital phase 0.5 (PI: Nowak), and have an approved observation of orbital phase 0.25 (PI: Hanke). Coupled with archival \chandra-\hetg observations of orbital phases near 0.75, we hope to develop a better understanding of how the ionized absorption evolves with orbital phase, which should allow us to improve our spectral modeling in the future.}.} Again, we have not found an inner radius for this line that exceeds $\approx 40~GM/c^2$. Given the variations of this line with presumed continuum, however, we are as of yet unable to use this line for more refined diagnostics such as estimates of black hole spin. Although we are unable as of yet to draw firm conclusions as to the best geometry and physical mechanisms to describe the hard state spectra of \cyg, these new joint \suzaku-\rxte data provide a stunning contrast to our prior results using solely \rxte data \citep{wilms:06a}. For the \rxte data alone we were able to describe the 3--125\,keV spectra with a variety of physically motivated Comptonization models and to describe correlations among the fit parameters. On the other hand, the simple exponentially cutoff, broken powerlaw models with a single, broad gaussian line described the data equally well, if not better \citep{wilms:06a}. When considering the 0.8--300\,keV \suzaku-\rxte data discussed here, this is no longer the case. We now require complex absorption at low energy, an asymmetric broad line plus narrow emission and absorption components in the Fe line region, and a complex continuum model. The catalog of \suzaku observations of BHC will continue to increase such that we observe a wider variety of BHC states and luminosities. Furthermore the sophistication and physical self-consistency of the spectral models will continue to improve. Together, they offer the promise of obtaining a better understanding of the physical processes occurring in these BHC systems.

1012.4801

1012

1012.4530.txt

We present a statistic for the detection of stochastic gravitational wave backgrounds (SGWBs) using radiometry with a network of multiple baselines. We also quantitatively compare the sensitivities of existing baselines and their network to SGWBs. We assess how the measurement accuracy of signal parameters, e.g., the sky position of a localized source, can improve when using a network of baselines, as compared to any of the single participating baselines. The search statistic itself is derived from the likelihood ratio of the cross correlation of the data across all possible baselines in a detector network and is optimal in Gaussian noise. Specifically, it is the likelihood ratio maximized over the strength of the SGWB and is called the maximized-likelihood ratio (MLR). One of the main advantages of using the MLR over past search strategies for inferring the presence or absence of a signal is that the former does not require the deconvolution of the cross correlation statistic. Therefore, it does not suffer from errors inherent to the deconvolution procedure and is especially useful for detecting weak sources. In the limit of a single baseline, it reduces to the detection statistic studied by \citet{Ballmer} and \citet{Mitra}. Unlike past studies, here the MLR statistic enables us to compare {\em quantitatively} the performances of a variety of baselines searching for a SGWB signal in (simulated) data. Although we use simulated noise and SGWB signals for making these comparisons, our method can be straightforwardly applied on real data. %CHECK: ({\color{cyan} need to address: how to apply on real search including confidence of detection}).

Just like the discoveries of the cosmic microwave background and pulsars in the electromagnetic spectrum, a discovery of unknown sources by Earth-based detectors such as LIGO and Virgo in the gravitational wave (GW) spectrum by serendipity is an interesting prospect. The LIGO Scientific Collaboration and the Virgo Collaboration are addressing it by searching for both transient signals, or ``bursts,'' and long-duration signals in the data from their detectors. Here, we focus on a subset of the latter type that can be modeled as a stochastic background. The search for an isotropic stochastic GW background has caught significant attention due to its cosmological significance. This primordial GW background is a direct probe of cosmological inflation \cite{Abbott:2009ws}. However, the astrophysical background, arising in the nearby Universe~\cite{Coward}, e.g., from an unresolved superposition of GW signals from multiple sources, such as low-mass x-ray binaries or, even, coalescing compact objects, is possibly much stronger than the primordial background and is anisotropic. A variety of data analysis strategies to search for an anisotropic GW background have been proposed and implemented in the past~\cite{AllenOttewill,cornish,kudohI,kudohII,taruyaIII,Cannon:2007br}. These searches are usually performed in two types of bases in the sky, namely, the pixel and spherical harmonic bases. Use of the radiometer technique for searching a GW background was proposed in Ref. \cite{LazzariniWeiss} and was implemented in the pixel basis on data from LIGO's fourth science run~\cite{S4Radiometer}. An elaborate study of this method, including the maximum-likelihood (ML) estimation of the true anisotropy of GW background by deconvolving the observed sky map, was presented in \citet{Mitra}. Even though the pixel-based search is promising and simpler to understand, it is not the best basis for probing sources with angular spreads greater than the angular resolution of the GW radiometer. The spherical harmonic basis is expected to be better suited for detecting such sources \cite{ThraneEtal}. Past attempts at probing the GW anisotropy in the spherical harmonic basis were essentially studies of the periodic modulation of the observed background in the detector baselines. Recently, a general ML formalism was developed to search for the GW anisotropies in any basis, including the spherical harmonic basis, using a network of ground-based GW interferometers \cite{ThraneEtal}. The pixel-based search is a specific application of this formalism. One of the main goals of this paper is to perform a thorough comparison of the expected performances of individual baselines and the whole network in detecting an astrophysical stochastic gravitational wave background (SGWB) and in estimating its parameters. The pixel basis is used for this study. %-- was yet to be done and that is one of the main goals of this paper. Even though a pixel-based search is optimal for a localized source, the resolution of the source is limited by the length of the radiometer baselines, the orientation of the detectors, and their individual sensitivities. % and can span tens of pixels (WHAT IS THE SIZE OF A PIXEL?!). Probing a stochastic GW background with %multiple observed values energy distributed across the pixelated sky demands a statistically meaningful integration of the energies received in every pixel. In order to accomplish this, we extend the maximized-likelihood ratio (MLR) statistic for a single baseline to incorporate a network of detectors or, equivalently, multiple baselines. The rest of the paper is devoted to studying the performance of individual GW detector baselines and the whole network by comparing different figures of merit for their performance, e.g., sensitivity, accuracy in localizing sources, sky coverage, and faithful extraction from the data of the sky distribution of a stochastic %true (simulated) background. The paper is organized as follows: In Sec.~\ref{sec:statistic}, we develop and study the efficiency of an optimal all-sky search statistic for anisotropic SGWBs that obviates the solving of the inverse problem, which may not always be well-posed. In Sec.~\ref{sec:performance}, we compare the performance of a network with that of its individual baselines using a variety of figures of merit. %CHECK2: ``numerical results of WHAT?'' The numerical results are also presented in this section. In Sec.~\ref{sec:conclusion}, we conclude by summarizing the implications of this work on ongoing SGWB searches and by highlighting future directions in GW radiometry. %The aim of this paper is to is to (1) find the maximum-likelihood (ML) statistic for detecting an astrophysical GW background (AGWB) with a single baseline, (2) extend it for multiple baselines, (3) compare their sensitivities to different sky positions for the LIGO-Virgo baselines, and (4) compare the accuracies with which a single-pixel source can be located with the separate baselines and their network. In the process, we establish a formalism in which the effectiveness of pixel- and spherical-harmonic-based deconvolution techniques for obtaining source sky-maps can be compared for different source distributions. %{\color{red} REMOVE:} {\color{magenta} Apart from an isotropic stochastic gravitational-wave (GW) background, predicted by slow-roll inflationary models as arising in the primordial universe, we also expect astrophysical backgrounds arising from a group of unresolved transient sources, e.g., coalescing binaries of compact objects, such as neutron stars and black holes \cite{Coward}. Such a background may contain anisotropic components. Moreover, ill-modeled astrophysical sources can also give rise to localized SGWB signals. The data analysis pipeline explored in Refs. \cite{Flanagan, AllenRomano} was optimal for searching isotropic SGWBs. By comparison, the pipelines explored in Refs. \cite{LazzariniWeiss, Ballmer, Mitra, ThraneEtal} addressed the problem of estimating the power distribution of a SGWB from the latter category that is anisotropic or localized. This work extends that body of work in the following ways. First, we explicitly obtain a detection statistic for optimally searching anisotropic SGWBs. An alternative statistic was implied in the literature and will be compared with our statistic here. Second, we compare sensitivities of different networks to ``single-pixel" sources as functions of sky declination and, for monochromatic sources, the source frequency. Third, we derive the parameter-space metric on the sky for searching localized SGWBs with this statistic. This allows us to compare different the resolving power of different detector networks. Finally, we assess the effect of systematic error owing to an inaccurate modeling of the spherical-harmonic modes of a distribution on the estimated parameters of that background. (Complete this list.) % %Past work on ASGWB focused more on estimating the parameters of such a background, such as the spherical harmonic modes of its power distribution across the sky. If the power in any single mode is loud, that parameter itself can be used as a detection statistic. It is, however, conceivable that there are ASGWBs that do not have significant power in any mode, but have significant power integrated across all modes. For such an ASGWB, this latter statistic is a more optimal one to employ for detecting it.} %

\label{sec:conclusion} The search for an anisotropic stochastic gravitational wave background plays an important role in present GW research. In addition to setting interesting upper limits on astrophysical and cosmological backgrounds, the simplicity of the concomitant analysis reveals invaluable knowledge about the coherent performance of the GW detector network. %CHECK: [{\color{red} Sanjit: I do not understand the last part of the preceding sentence?}] So far, detailed analysis strategies have been developed to search for anisotropic background in pixel and spherical harmonic spaces, and a general maximum-likelihood-based framework has been established to search in any convenient basis. The spherical harmonic search has been demonstrated using a network of detectors~\cite{ThraneEtal}. In this paper, for the first time, we numerically implement the directed radiometer search, including deconvolution, for a network of detectors. These methods, in the past, were focused primarily on showing that the observed map is consistent with Gaussian noise or in estimating sky maps. The latter required the inversion of the convolution equation, which itself assumed the network of detectors to be nondegenerate. Neither of these methods may work in the presence of excessive noise and weak signal. Most importantly, a statistically meaningful, all-sky combined statistic, in the form of an optimal ``detection statistic,'' was needed in order to make precise statements about the presence or absence of a given background model in a map. Here, we proposed a MLR statistic, which yields a single number when computed on the dirty or the clean map and can be used as a detection statistic. By computing the MLR statistic for a couple of toy models of the background, we observe that the detection statistic is much larger than the noise-only case, even in the presence of weak signals that are barely visible in dirty or clean maps. We corroborated these statements with results obtained from extensive Monte Carlo simulations of a diffuse background of known shape in an ensemble of noise realizations. However, a more detailed study using signals from a variety of background models is surely worth pursuing in order to determine how accurate the templates need to be in order to extract meaningful information from weak backgrounds. %CHECK: The numbers are promising for the basic exercise we performed here. We also compared the performance of individual baselines and the whole network for the directed radiometer search using different figures of merit. %CHECK: Note that, the performance analysis of a network of GW detector presented here is relatively straightforward compared to other GW searches~\cite{Bose00,Pai01,WenChen04,Hayama07,Dhurandhar08, Gursel89} and, hence, can be useful in drawing insights about a network. Evaluating the performance of a network of GW detectors in SGWB searches is relatively straightforward compared to other GW signal searches~\cite{Bose00,Pai01,WenChen04,Hayama07,Dhurandhar08, Gursel89}. This exercise was useful in drawing insights about the characteristics of a network that are particularly helpful in boosting its performance. %CHECK: Not sure how this is related to the first half of the sentence: ``and, hence, can be useful in drawing insights about a network.'' Our overall observation, not surprisingly, is that the network improves performance in mainly three ways, namely, (1)~by increasing the sensitivity by observing each direction a greater number of times, (2)~by observing the sky more uniformly, and (3)~by probing each direction on the sky with additional detectors on the globe. The latter two enhancements lead to better localization of pointlike sources. This can be understood via the behavior of the Fisher information matrix: More detectors reduce its degeneracy and improve the well-posedness of the inverse problem. This, in turn, leads to a more accurate production of clean maps. Another question worth addressing in the future is about how closely spaced must the templates be on the parameter space to maximize the chances of detection with available computational resources. Indeed, the proposal for templated searches for SGWB signals is not new to this paper. For example, it has been addressed earlier in the context of isotropic searches (see Ref. \cite{Bose:2005fm} and the references therein). Reference \cite{Bose:2005fm} also introduced a metric on the parameter space of those signals so as to enable an experimenter to infer what the principle axes are on that space and how fine a template bank one can afford based on the computational resources available. A similar study can be carried out for finding a more optimal spacing of templates for directed searches than the one used here. Whereas results presented here were derived for Gaussian noise, the codes used can be applied to real data as well. Indeed, the performance of the proposed statistic in real data sets from the LIGO and Virgo detectors can be determined through hardware injections that were done in the recent science runs, such as the ones described in Ref. \cite{Bose:2003nb}, and supplementing them with multiple software injections to improve the statistics. The expected improvement of network sensitivity over individual baselines, as demonstrated here, merits the investment required for extending the current single-baseline analysis efforts \cite{Mitra,Ballmer} to a multibaseline one. This conclusion is strengthened by the fact that adding a detector to a baseline can potentially mitigate the contribution of cross correlated environmental noise that affects only one of the three resulting baselines. Including V1, which is on a different continental plate than the H1L1 baseline, can serve this purpose. Employing a null-stream statistic~\cite{Chatterji06,Gursel89} to complement the detection statistic might also help in discriminating against such noise.

1012.4530

1012

1012.2997_arXiv.txt

Luminous Blue Variables (LBVs) are massive post-Main Sequence stars that are experiencing a highly unstable phase of evolution that is characterised by dramatic photometric and spectroscopic variability and heavy mass loss. They have been the subject of much recent interest given the twin possibilities that their high mass-loss rates -- particularly during transient outbursts -- may be essential for the formation of H-depleted Wolf Rayet stars (e.g. Smith \& Owocki 2006) and that they may be the immediate precursors of a subset of highly luminous Type II supernovae (e.g. Gal-Yam \& Leonard 2009). Historically, their rarity (e.g. Clark et al. 2005) has meant that their properties -- particularly regarding their characteristic outbursts and eruptions (duration, duty cycle, associated mass loss rate and underlying physical cause) -- have remained poorly understood. However, recent narrow- and broad-band infra-red surveys of the Galactic Plane have revealed a large number of new LBVs candidates (Clark et al., 2003, Gvaramadze et al. 2010, Mauerhan et al. 2010, Wachter et al. 2010) and it is hoped that studies of an expanded sample size will help elucidate the nature of the LBV phenomenon and its role in massive stellar evolution. However, given their location in the Galactic plane, observations of these stars must be undertaken in the (near)-IR due to significant line of sight extinction. In this contribution we preview the results of a long term spectroscopic and photometric campaign of recently identified candidate LBVs, supplemented with tailored model atmosphere analysis utilising the CMFGEN code (Hillier \& Miller 1998); a full description of this program will be presented in Clark et al. (in prep.).

While various lines of evidence suggest an important role for LBVs in the late evolutionary stages of massive stars -- and by extention their death in SNe and the nature of the post-SNe relativistic remnants -- the properties of this phase are still poorly understood, in large part due to the rarity of such stars. However, the recent identification of large numbers of new candidates within the Galactic plane and the viability of studying them via concerted spectroscopic and photometric monitoring supplemented with tailored non-LTE model atmosphere analysis will allow these issues to be directly addressed. Indeed, such work will greatly benefit from near-IR surveys such as VISTA/VVV and the advent of 1-2m class robotic facilities such as the Faulkes Telescopes. Likewise, the availability of multiplexing spectrographs and transient surveys such as PanSTARRS will permit similar studies in external galaxies over a range of metalicities. When combined with radiative transfer modeling of the spatially resolved gaseous \& dusty ejecta associated with large numbers of galactic LBVs and which encodes their past mass loss histories, these programs have the potential to advance studies of this transient and violent phase of stellar evolution over the coming years. \begin{figure}[h] \begin{minipage}{8cm} \centering \includegraphics[width=6cm,angle=270]{JS_Clark1_fig1.ps} \end{minipage} \hfill \begin{minipage}{8cm} \centering \includegraphics[width=6cm,angle=270]{JS_Clark1_fig2.ps} \end{minipage} \caption{ Long term JHK lightcurves of G24.73+0.69 (left panel; Clark et al. in prep) and AFGL2298 (right panel; Clark et al. 2009). Times of spectral observations are indicated by arrows.} \end{figure} \begin{figure}[h] \centering \includegraphics[width=17cm]{JS_Clark1_fig3.ps} \caption{Montage of K band spectra of galactic (candidate) LBVs demonstrating their diverse morphologies. For comparison the spectra of the Yellow Hypergiant IRC +10 420 and the WN8 Quintuplet member LHO 158 are also presented. For reasons of brevity the references to the origin of spectra and associated analyses have been omitted but are given in Clark et al. (in prep.). } \end{figure}

1012.2997

1012

1012.0730_arXiv.txt

We compare the three-dimensional gas temperature distributions obtained by a dedicated radiative transfer and photoionisation code, {\sc MOCASSIN}, against those obtained by the recently-developed Smooth Particle Hydrodynamics (SPH) plus ionisation code {\sc iVINE} for snapshots of an hydrodynamical simulation of a turbulent interstellar medium (ISM) irradiated by a nearby O star. Our tests demonstrate that the global ionisation properties of the region are correctly reproduced by {\sc iVINE}, hence validating further application of this code to the study of feedback in star forming regions. However we highlight potentially important discrepancies in the detailed temperature distribution. In particular we show that in the case of highly inhomogenous density distributions the commonly employed on-the-spot (OTS) approximation yields unrealistically sharp shadow regions which can affect the dynamical evolution of the system. We implement a simple strategy to include the effects of the diffuse field in future calculations, which makes use of physically motivated temperature calibrations of the diffuse-field dominated regions and can be readily applied to similar codes. We find that while the global qualitative behaviour of the system is captured by simulations with the OTS approximation, the inclusion of the diffuse field in {\sc iVINE} (called {\sc DiVINE}) results in a stronger confinement of the cold gas, leading to denser and less coherent structures. This in turn leads to earlier triggering of star formation. We confirm that turbulence is being driven in simulations that include the diffuse field, but the efficiency is slightly lower than in simulations that use the OTS approximation.

Ionising radiation from OB stars influences the surrounding interstellar medium (ISM) on parsec scales. As the gas surrounding a high mass star is heated, it expands forming an HII region. The consequence of this expansion is twofold, on the one hand gas is removed from the centre of the potential preventing it from further gravitational collapse. On the other hand gas is swept up and compressed beyond the ionisation front producing high density regions that may thus be susceptible to gravitational collapse (i.e. the ``collect and collapse'' model, Elmegreen et al 1995). Furthermore, pre-existing, marginally gravitationally stable, clouds may also be driven to collapse by the advancing ionisation front (i.e. ``radiation-driven implosion'', Bertoldi 1989, Kessel-Deynet \& Burkert 2003, Gritschneder et al 2009a). Finally, ionisation feedback is also thought to be a driver for small scale turbulence in a cloud (Gritschneder et al 2009b). The net effect of photoionisation feedback on the global star formation efficiency is still, however, under debate. \begin{figure*} \begin{center} \includegraphics[width=18cm]{fig1_comp_div.png} \caption{Surface density projected in the z-direction of the t~=9~kyr (left), 250~kyr (middle) and ~500~kyr (right) snapshots of G09b's turbulent ISM simulation. The upper panels show the evolution of the gas without diffuse field, the lower panels show the same simulation performed with the simplified diffuse field implementation discussed in Section~4. The white boundaries on the left-hand panels mark the regions compared in Table~1. The SPH particles were mapped onto a 128$^3$ regular Cartesian grid.} \end{center} \end{figure*} While the importance of studying the photoionisation process as part of hydrodynamical star formation simulations has long been widely recognised, until very recently, due to the complexity and computational demand of the problem, the evolution of ionised gas regions had only been studied in rather idealised systems (e.g. Yorke et al. 1989; Garcia-Segura \& Franco 1996), with simulations often lacking resolution and dimensions. Fortunately, the situation in the latest years has been rapidly improving, with more sophisticated implementations of ionised radiation in grid-based codes presented by (e.g.) Mellema et al (2006), Krumholz et al (2007) and Peters et al (2010). Kessel-Deynet \& Burkert (2000) were the first to introduce an ionisation algorithm into a Smoothed Particle Hydrodynamical (SPH) code to study radiation-driven implosion as a possible trigger of star formation. Later, Dale, Ercolano \& Clarke (2007) presented a much simplified, but fast, algorithm to consider photoionisation within complex SPH simulations. When compared to grid codes, which are based on the solution of the Eulerian form of the same equations, much higher resolution of very complex flows can be achieved. Since Dale et al (2007) a number of other ionising radiation implementations have been developed for SPH codes, including Pawlik \& Schaye (2008), Altay et al. (2008), Bisbas et al (2009) and very recently Gritschneder et al (2009a). However high resolution SPH simulations are very computationally expensive, and even in the current era of parallel computing, an exact solution of the radiative transfer (RT) and photoionisation (PI) problem in three dimensions within SPH calculations is still prohibitive. Necessarily, all the algorithms mentioned above employ an extremely simplified approach to RT and PI. The consequences of such simplification on the conclusions drawn from the simulations need to be investigated. In Dale et al (2007), we performed the only such verification to date against a fully three-dimensional radiative transfer and photoionisation code ({\sc MOCASSIN}, Ercolano et al 2003, 2005 and 2008) for complex density fields obtained from the SPH calculations. We showed in that case that the agreement on the ionised fractions in high density regions was very good, but low density regions were poorly represented by the ionisation + SPH code. In this paper we take a similar approach to test the more recent algorithms developed in {\sc iVINE} (Gritschneder et al. 2009a). We find excellent agreement between the codes for the global ionisation fractions, however we highlight discrepancies in the temperature distribution, particularly in shadow regions that are dominated by the diffuse field, which is not accounted for in iVINE. We test the consequences of the omission of the diffuse field both on the hydrodynamical evolution of the structure and on the evolution of the turbulence spectrum using a simple approach that allows for a more realistic, but still efficient modelling strategy of the shadow regions. In Section~2 we briefly describe the {\sc iVINE} and {\sc MOCASSIN} codes and the comparison strategy. In Section~3 we show the results of this comparison. In Section~4 we discuss a simple approach to qualitatively include diffuse field effects in {\sc iVINE} and show the results from these further tests in Section~5. Section~6 contains a brief summary and future directions.

We have presented a detailed comparison of the ionisation and temperature structure for a turbulent ISM simulation performed with the SPH+ionisation code {\sc iVINE} (Gritschneder et al 2009a,b, 2010) against the solution obtained with the photoionisation code {\sc MOCASSIN} for snapshots of the density distribution. {\sc iVINE} treats hydrodynamics, gravitational forces and ionisation simultaneously. The ionisation is calculated by making use of the 'on-the-spot' (OTS) approximation. The {\sc MOCASSIN} code (Ercolano et al 2003, 2005, 2008a) is fully three-dimensional and includes an exact treatment for the frequency resolved transfer of both the stellar (direct) and diffuse radiation fields. {\sc MOCASSIN} includes all the microphysical processes that dominate the thermal and ionisation balance of the ionised gas, providing realistic temperature and ionisation distributions. Our tests show that {\sc iVINE} and {\sc MOCASSIN} agree very well on the global properties of the region (i.e. total ionised mass fraction and location of the main ionisation front), but we note discrepancies in the temperature structure, particularly in the shadow regions. These tend to be cold and neutral in the {\sc iVINE} plane-parallel stellar-field-only prescription , while {\sc MOCASSIN} obtains a range of ionisation levels and temperatures that can be very crudely described as a function of density. We have developed a computationally inexpensive strategy to include the thermal effects of the diffuse field, as well as accounting for environmental variables, such as gas metallicity and stellar spectra hardness. The method relies on the identification of the shadow region via simple criteria and application of a temperature parameterisation that was obtained a-priori using {\sc MOCASSIN}. The method can be readily extended to other hydrodynamical codes (both SPH and grid-based). We evaluate the effects of diffuse fields by comparing runs with the standard {\sc iVINE} and the diffuse field implementation {\sc DiVINE}. In agreement with previous studies (Raga et al 2009), we find that the overall qualitative behaviour of the system (i.e. the formation of what appear to be pillar like structures) is similar in the two runs. Nevertheless our models demonstrate that the diffuse field has important quantitative effects on the hydrodynamical evolution of the irradiated ISM. In particular we note that {\sc DiVINE} predicts denser and less coherent structures which are much less attached to the parental cloud. This is due to the higher compression of the cold structures by the diffusively heated material inside the pillars trunks. Triggered star formation is promoted by this effect to a much earlier time. The compression also affects the turbulence spectrum of the system. We confirm the driving of turbulence by the ionising radiation, but with a slightly reduced efficiency compared to previous calculations with {\sc iVINE} using the OTS approximation.

1012.0730

1012

1012.0506_arXiv.txt

Since the detection of the asymptotic properties of the dipole gravity modes in the Sun, the quest to find individual gravity modes has continued. An extensive and deeper analysis of 14 years of continuous GOLF/SoHO observational data, unveils the presence of a pattern of peaks that could be interpreted as individual dipole gravity modes in the frequency range between 60 and 140 microHz, with amplitudes compatible with the latest theoretical predictions. By collapsing the power spectrum we have obtained a quite constant splitting for these patterns in comparison to regions where no g modes were expected. Moreover, the same technique applied to simultaneous VIRGO/SoHO data unveils some common signals between the power spectra of both instruments. Thus, we are able to identify and characterize individual g modes with their central frequencies, amplitudes and splittings allowing to do seismic inversions of the rotation profile inside the solar core. These results open a new ligh t on the physics and dynamics of the solar deep core.

Little progress has been done during the last few years on the structure [1,2,3] and Dynamics [4,5,6,7] of the solar core even after the detection of the asymptotic spacing of the dipole gravity modes [8,9] Indeed no general consensus has been obtained for the detection of individual g modes yet [10]. This is because the increasing convective background level towards the lower frequencies (e.g. [11]), combined with very small amplitudes of those modes (several mm/s in the case of g modes, [12]) are the limiting factors for their detections. In the case of the low-degree, low-frequency p modes, accurate measurements are hardly obtained below 1~mHz [13,14,15]. This situation might change in the near future when we have a few years of data coming from the new instrumentation available, such as PICARD [16], and the very promising HMI and AIA aboard SDO [17] or the new projected instrumentation (e.g. GOLF-NG [18]). Fortunately, time goes by and the signal-to-noise ratio of the Global Oscillations at Low Frequency (GOLF) instrument [19] and the Sun Photometers (SPM) on the Variability of IRradiance and Global Oscillations [20] aboard the Solar and Heliospheric Observatory (SoHO) mission increases. In this work we uncover the presence of peaks in the power spectral density (PSD) that could be individual $\ell$=1 modes. These peaks are regularly spaced in period in the positions determined by the asymptotic periodicity measured by [8] Moreover, these peaks present a regular pattern in frequency that could be the signature of the rotational splitting. Thus this work study potential candidates that might merit further investigation, a step forward of what it has been done during the last few years (e.g. [21,22,23]).

In this paper we have been able to identify the individual peaks generating the periodic signal found by [8] and interpreted as the asymptotic periodicity of the dipole gravity modes. The analysis of the collapsograms gave five possible detections in the frequency range between 60 and 140 $\mu$Hz with a rotational splittings in the range 850 to 950 nHz. When these candidate modes are used in the inversion of the rotational profile, we see that the rotation in the deep core reaches values close to 4000 nHz. However, more work should be done to better characterize these g-mode candidates as well as the determination of the error bars that are very important for the inversions. \ack SoHO is a space mission of international cooperation between ESA and NASA. DS acknowledges the support of the grant PNAyA 2007-62650 from the Spanish National Research Plan, J.B. the support through the ANR SIROCO, and R.A.G and S.T. thanks the support form the CNES and the French Stellar Physics National Plan (PNPS). NCAR is supported by the National~Science~Foundation.

1012.0506

1012

1012.1500_arXiv.txt

The effects of initially uniform magnetic fields on the formation and evolution of dense pillars and cometary globules at the boundaries of H~\textsc{II} regions are investigated using 3D radiation-magnetohydrodynamics simulations. It is shown, in agreement with previous work, that a strong initial magnetic field is required to significantly alter the non-magnetised dynamics because the energy input from photoionisation is so large that it remains the dominant driver of the dynamics in most situations. Additionally it is found that for weak and medium field strengths an initially perpendicular field is swept into alignment with the pillar during its dynamical evolution, matching magnetic field observations of the `Pillars of Creation' in M16 and also some cometary globules. A strong perpendicular magnetic field remains in its initial configuration and also confines the photoevaporation flow into a bar-shaped dense ionised ribbon which partially shields the ionisation front and would be readily observable in recombination lines. A simple analytic model is presented to explain the properties of this bright linear structure. These results show that magnetic field strengths in star-forming regions can in principle be significantly constrained by the morphology of structures which form at the borders of H~\textsc{II} regions.

\label{sec:introduction} The study of elephant trunks, pillars and globules, found at the borders of H~\textsc{II} regions around massive stars, has received significant attention in recent years, both from an observational perspective and in theoretical and computational models. The famous `Pillars of Creation' in M16 were observed at optical wavelengths with HST~\citep{HesScoSanEA96}, in the IR~\citep{IndRobWhiEA07,SugWatTamEA07}, and in sub-mm/radio~\citep[e.g.][]{Pou98,WhiNelHolEA99}, showing that these are dynamic structures with ongoing star formation which may or may not have been triggered by the radiation which has shaped the pillars. \citet{SmiPovWhiEA10} provide strong evidence that pillars in the Carina Nebula are significant sites of sequential star formation propagating away from the older star clusters in this region, building on previous observations of synchronised star formation around the periphery of the nebula by~\citet{SmiBro07}. On smaller scales, studies of T-Tauri star ages in the Orion Nebula~\citep{LeeCheZhaEA05,LeeChe07} show decreasing stellar ages moving away from massive stars and towards bright-rimmed clouds at the H~\textsc{II} region/molecular cloud interface, again strongly suggesting at least sequential and possibly triggered star formation. The clear relationship between pillars/globules and second generation star formation around OB associations, and the question of the extent to which this star formation is triggered, provides strong motivation to understand the formation and evolution of these structures. In a previous paper~\citep[][hereafter ML10]{MacLim10} we investigated the formation and evolution of dense pillars of gas and dust -- elephant trunks -- on the boundaries of H~\textsc{II} regions using 3D hydrodynamical simulations including photoionising radiative transfer (R-HD). It was found that shadowing of ionising radiation by an inhomogeneous density field naturally forms elephant trunks without the assistance of self-gravity, or of ionisation front and cooling instabilities. A combination of radiation-driven implosion~\citep[RDI;][]{Ber89} and acceleration due to the rocket effect~\citep{OorSpi55} produce elongated structures: RDI compresses neutral gas until pressure equilibrium with ionised gas is achieved; the rocket effect accelerates gas away from the radiation source producing dynamic elongated structures with lifetimes of a few hundred kyr (depending on clump masses/densities). Strong neutral gas cooling was found to enhance this formation mechanism, producing denser and longer lived pillar-like structures compared to models with weak cooling. Models such as these for the formation of bright-rimmed clouds, globules and pillars have been considered for many years~\citep[e.g.][]{Pot58, Mar70, BodTenYor79, SanWhiKle82}; much of this work is summarised by~\citet{Yor86}. The RDI of a photoionised clump and its subsequent acceleration and evolution was calculated analytically~\citep{Ber89,BerMcK90} and subsequently numerically by~\citet{LeFLaz94}. \citet{WilWarWhi01} considered a range of axisymmetric models showing that multiple scenarios could form long-lived pillar-like structures. It has been shown \citep{KesBur03, MiaWhiNelEA06, PouKanRyuEA07, BisWhiWueEA10} that RDI of single clumps can generate cometary globules and trigger gravitational collapse, but it is more difficult to form pillars like those in M16 because gas must accumulate to a high density in the shadowed tail region. \citet{LimMel03} showed how photoionisation of multiple clumps which partially shadow each other leads to dense gas accumulating in shadowed tail regions; it was suggested by~\citet{PouKanRyuEA07} that multiple clumps are required to quantitatively match observations of the pillars in M16. Recent models~\citep[][ML10]{MelArtHenEA06, RagHenVasEA09, LorRagEsq09, GriNaaWalEA09, GriBurNaaEA10} showed that elephant trunks can form quite naturally from the photoionisation of a clumpy density field under a range of initial conditions. \citet{GriBurNaaEA10} extended our previous work (ML10) which used static initial conditions by considering the photoionisation of dynamic density fields generated by isothermal decaying turbulence, measuring the formation of pillar-like structures as a function of turbulent Mach number. Filamentary structure in an apparently helical geometry was found in the shadowed tail regions of a number of elephant trunks\ \citep{CarKriGah98, CarGahKri02, CarGahKri03}, and it was suggested this could arise due to twisting of magnetic field lines. Velocity profiles for some trunks were shown to be consistent with solid-body rotation~\citep{GahCarJohEA06}, again consistent with a magnetic origin of the observed structure, although the actual magnetic field orientation and strength has not yet been measured for these structures. The magnetic field in cometary globules has been measured by optical polarimetry~\citep{SriBhaRaj96,Bha99,BhaMahMan04}, showing in two cases a field orientation along the head--tail axis of the globule, but in one case~\citep{Bha99} a perpendicular field was found. \citet{SugWatTamEA07} used near-IR polarimetry of background stars to measure the magnetic field in M16, finding an ordered large-scale field in the H~\textsc{II} region, but within the pillars the field is aligned with the pillar axes and significantly misaligned with the ambient field by $\theta \sim 30$--$40^{\circ}$. They suggest this should constrain the magnetic field strength since it has not been strong enough to resist reorientation during the formation and evolution of the pillars, a suggestion we explore in more detail in this work. Theoretical calculations of the effects of magnetic fields on the expansion of H~\textsc{II} regions were first considered by~\citet{Las66b}. Using analytic calculations \citet{RedWilDysEA98} studied ionisation fronts with a plane-parallel magnetic field in the plane of the ionisation front. They found that the velocity separation between R-type and D-type solutions decreases with increasing field strength, and that D-critical ionisation fronts also advance into the neutral gas more rapidly for increasing field strength. In~\citet{WilDysHar00} jump conditions for ionisation fronts with oblique magnetic fields were presented, together with 1D numerical models showing how an ionisation front could progress from fast-R-type through fast-D-type, slow-R-type, and finally to slow-D-type. These extra modes are allowed because fast and slow shocks detach from the ionisation front at different times and propagate into the neutral medium. This work was extended by~\citet{WilDys01}, who calculated the internal structures of stationary 1D ionisation fronts. They found that oblique fields could produce significant transverse velocities and regions where the transverse field component is significantly modified. \citet{Wil07} used 2D slab-symmetric radiation-magnetohydrodynamics (R-MHD) simulations to study the photoevaporation flows from magnetised globules. In these models clumps with an initial density of $n_\mathrm{H}=2\times10^5\,\mathrm{cm}^{-3}$ were placed in an ambient medium $10\times$ less dense. For simplicity a uniform magnetic field with various strengths and orientations was used. Plane-parallel radiation was assumed, with a thermal model where the temperature relaxed to a value between $100\,$K and $10\,000\,$K according to its ionisation fraction. It was found that for a weak field the ionised gas pressure dominates the dynamics and the field was swept into a configuration where it was parallel to the column of neutral gas behind the dense clump. For a sufficiently strong field, however, the field determined the dynamics and made the flow almost one-dimensional along field lines. The first 3D R-MHD calculation including non-equilibrium photoionising radiative transfer was performed by~\citet*{KruStoGar07}. They used the \textsc{Athena} code with a new ray-tracing module to simulate the expansion of an H~\textsc{II} region around a point source in a uniform magnetic field. The overall expansion is now axisymmetric rather than spherically symmetric~\citep[cf.][]{Las66a} with a dense shell forming in directions parallel to the field, and a possibly numerical instability developing for expansion perpendicular to the field. Using similar methods,~\citet{HenArtDeCEA09} model the photoionisation of a dense clump of gas in 3D with an initially uniform magnetic field. They use the photon-conserving C$^2$-ray ray-tracing algorithm~\citep{MelIliAlvEA06} which allows large timesteps to be taken without loss of accuracy. They found that the evolution of a photoionised globule can be significantly altered by the presence of a strong field, and in some cases a recombining shell forms at the termination shock of the photoevaporation flow when it is confined by a transverse field. The implosion phase is strongly asymmetric since the clump compresses much more readily along field lines than across them. These authors introduce a detailed thermal model to model the dynamics as realistically as possible for conditions in the Orion Nebula, finding that X-ray heating keeps the neutral gas temperature at $T\gtrsim50\,$K, but the cooling in neutral gas is somewhat stronger than that considered in ML10. In this work we add magnetic fields of various strengths and orientations to some of the models considered in ML10 to study their effects on the dynamics of the dense neutral gas. The simulation code is described in \S\ref{sec:numerics}: \S\ref{ssec:MHD} describes the MHD dynamics algorithm; \S\ref{ssec:MP} reviews the microphysical processes included; \S\ref{ssec:KSG07} presents a brief comparison of results obtained with our code to the test problem described by~\citet{KruStoGar07}. The 3D simulations presented here are introduced in \S\ref{sec:simulations} and our results presented in \S\ref{sec:results}. \S\ref{ssec:projections} shows the evolution of the projected gas density and magnetic field orientation; \S\ref{ssec:emission} shows the emission from ionised gas; \S\ref{ssec:R8R8a} evaluates boundary effects in strong field simulations. A calculation explaining the presence of a bright ionised linear ridge/ribbon in the strong field simulations is presented in \S\ref{sec:BarFormation}. These results are compared to observations of H~\textsc{II} regions and their magnetic fields in \S\ref{sec:discussion}, where we also discuss our work in the context of other recent computational research. Our conclusions are presented in \S\ref{sec:conclusions}, and some technical details and tests for the simulations are given in the appendices.

\label{sec:conclusions} We have performed a series of R-MHD simulations of the photoionisation of dense clumps of gas and their evolution from pillar-like to cometary globule-like structures. Our results for the emissivity of ionised gas agree very well with those of \citet{HenArtDeCEA09}, showing that a dense, ionised, bar-shaped region standing off from the ionisation front is a generic feature of strongly magnetised photoionisation in a clumpy medium for a perpendicular field orientation. This ridge can be as bright as, or even brighter than, the photoionisation front when observed in recombination radiation (e.g.~$\mathrm{H}\alpha$) and its presence or absence can be used as a diagnostic of the strength of any large scale magnetic field which may be present. Bright ridges or ribbons are observed in some H~\textsc{II} regions~\citep[e.g.][]{BraGreChuEA00, BohTapRotEA04, SmiBalWal10}, although they are not common. An overdense ridge could also be produced by ram-pressure confinement, and more detailed modelling is required to find observational signatures which could distinguish these different confinement mechanisms. Comparing to observations of M16~\citep{HesScoSanEA96} there is no such ribbon or ridge, suggesting the ambient field measured by~\citet{SugWatTamEA07} is not dominant in the ionised gas. This conclusion is strengthened when we consider the magnetic field orientation observed in M16~\citep[][fig.~9]{SugWatTamEA07}. The results presented here show that both RDI and acceleration of clumps by the rocket effect tend to align the magnetic field in dense neutral gas with the radiation propagation direction. In our models a field configuration similar to the observed one is clearly seen when the initial field strength is $\vert\mathbf{B}\vert\simeq20\,\mu$G; the simulation with $\vert\mathbf{B}\vert\simeq50\,\mu$G is consistent with observations, and the simulation with $\vert\mathbf{B}\vert\simeq160\,\mu$G is not consistent. Our simulations thus suggest an ambient field strength of $\vert\mathbf{B}\vert\lesssim 50\,\mu$G around the M16 pillars. The morphology of the structures which develop due to RDI and the rocket effect is also affected by a strong magnetic field, partly due to shielding by the dense ionised ridge and partly by the effect of the field within the pillar or globule. Inspection of $\mathrm{H}\alpha$ images of elephant trunks and globules in the literature~\citep[e.g.][]{HesScoSanEA96, SmiBalWal10} suggests that the features seen in the strong field simulations are not common, although the uniform initial field configurations considered here are certainly somewhat artificial. Additionally many H~\textsc{II} regions have significantly higher gas pressure than that modelled in our simulations, in which case the magnetic field must also be correspondingly stronger to dominate the dynamics. Comparing these results with our earlier simulations in ML10, the larger simulation domains used here show that the pillar-like structures which form will ultimately evolve to cometary structures in the absence of dense gas further from the star. The lifetimes of pillars in our models are $t \lesssim 500\,$kyr, although this depends significantly on the initial mass and concentration (and presumably velocity, cf.~\citealt{GriBurNaaEA10}) of the dense gas clumps. Finally we emphasise, in agreement with previous authors~\citep{Wil07,KruStoGar07,HenArtDeCEA09}, that a strong magnetic field has a very significant influence on the dynamics of the photoionisation process, and many of these effects should be easily observable. Given the difficulty of measuring the full 3D magnetic field in the ISM, comparison to detailed numerical simulations such as these offers an indirect means to constrain the field strength and orientation in and around H~\textsc{II} regions. \vspace{-0.5cm}

1012.1500

1012

1012.3505_arXiv.txt

We report on the analysis of $\sim$ 22,000 M dwarfs using a statistical parallax method. This technique employs a maximum--likelihood formulation to simultaneously solve for the absolute magnitude, velocity ellipsoid parameters and reflex solar motion of a homogeneous stellar sample, and has previously been applied to Galactic RR Lyrae and Cepheid populations and to the Palomar/Michigan State University (PMSU) survey of nearby low-mass stars. We analyze subsamples of the most recent spectroscopic catalog of M dwarfs in the Sloan Digital Sky Survey (SDSS) to determine absolute magnitudes and kinematic properties as a function of spectral type, color, chromospheric activity and metallicity. We find new, independent spectral type-absolute magnitude relations, and color-absolute magnitude relations in the SDSS filters, and compare to those found from other methods. Active stars have brighter absolute magnitudes and lower metallicity stars have fainter absolute magnitudes for stars of type M0-M4. Our kinematic analysis confirms previous results for the solar motion and velocity dispersions, with more distant stars possessing larger peculiar motions, and chromospherically active (younger) stars having smaller velocity dispersions than their inactive counterparts. We find some evidence for systematic differences in the mean $U$ and $W$ velocities of samples subdivided by color.

M dwarfs are the most numerous stellar population in the Milky Way \citep{2010AJ....139.2679B}. Surveys such as the Sloan Digital Sky Survey \citep[SDSS;][]{2000AJ....120.1579Y} and the Two Micron All Sky Survey \citep[2MASS;][]{2006AJ....131.1163S} have led to photometric and (in the case of SDSS) spectroscopic catalogs containing large numbers of low mass stars. The largest spectroscopic database of M dwarfs \citep[][]{west10} compiles spectral types, colors, proper motions, radial velocities and chromospheric activity estimates (as traced by Balmer series emission) for more than 70,000 stars from the most recent SDSS data release. In order to make widespread use of these new measurements of the low--mass stellar population, with a focus on Galactic structure and kinematics, it is necessary to have good distance estimates. Since only very nearby M dwarfs have measured trigonometric parallaxes, photometric and spectroscopic parallax relations have typically been employed to obtain distances based on a star's color or spectral type respectively \citep{2002AJ....123.3409H, 2005PASP..117..706W, 2007ApJ...662..413K, 2008ApJ...673..864J}. Because SDSS photometry saturates at $m \sim$ 15, it is particularly difficult to anchor the photometric parallax relations in the SDSS filters with measured trigonometric parallax stars. The classical statistical parallax method is a way to determine the absolute magnitude of a homogeneous set of stars. Statistical parallax analysis seeks to determine the distance scale that provides the best match between the measured proper motions and radial velocities of a given stellar population, returning an estimate of the average absolute magnitude of the population, and the kinematic properties including the reflex solar motion and the velocity ellipsoid (velocity dispersions along three principal axes). Previous discussions of the statistical parallax method can be found in \cite{1971MNRAS.151..231C}, \cite{1983veas.book.....M} and \cite{1998ApJ...506..259P}. Our particular formulation has been used to study RR Lyraes \citep {1986ApJ...302..626H, 1986MNRAS.220..413S, 1996AJ....112.2110L,1998A&A...330..515F,1998ApJ...506..259P}, Cepheids \citep{1991ApJ...378..708W}, and the nearby low-mass stars from the PMSU survey \citep{1996AJ....112.2799H}.

1012.3505

1012

1012.3169_arXiv.txt

Active galactic nuclei (AGN) have been observed to vary stochastically with $10-20\%$ rms amplitudes over a range of optical wavelengths where the emission arises in an accretion disk. Since the accretion disk is unlikely to vary coherently, local fluctuations may be significantly larger than the global rms variability. We investigate toy models of quasar accretion disks consisting of a number of regions, $n$, whose temperatures vary independently with an amplitude of $\sigma_T$ in dex. Models with large fluctuations ($\sigma_T=0.35-0.50$) in $10^{2-3}$ independently fluctuating zones for every factor of two in radius can explain the observed discrepancy between thin accretion disk sizes inferred from microlensing events and optical luminosity while matching the observed optical variability. For the same range of $\sigma_T$, inhomogeneous disk spectra provide excellent fits to the \emph{HST} quasar composite without invoking global Compton scattering atmospheres to explain the high levels of observed UV emission. Simulated microlensing light curves for the Einstein cross from our time-varying toy models are well fit using a time-steady power-law temperature disk, and produce magnification light curves that are consistent with current microlensing observations. Deviations due to the inhomogeneous, time-dependent disk structure should occur above the $1\%$ level in the light curves, detectable in future microlensing observations with millimag sensitivity.

Optical emission from active galactic nuclei (AGN) is thought to be due to thermal emission from a standard thin accretion disk \citep{shaksun1973,novthorne}. This model is successful in explaining many observations of black hole accretion disks, such as the peak frequencies of the thermal spectra of X-ray binaries and AGN. It is also used to fit continuum spectra in the X-ray to measure black hole spin \citep[e.g.,][]{shafee2006}, and to infer accretion disk size from gravitational microlensing; both in absolute terms \citep[e.g.,][]{morganetal2010} and as a function of wavelength \citep{anguitaetal2008,poindexteretal2008,eigenbrodetal2008,morganetal2010}. Attempts to fit continuum AGN spectra with thin disk models have met with mixed success \citep{blaesetal2001,davisetal2007}. While multi-temperature blackbody spectra generally fit well in the optical, at UV wavelengths quasar spectra are brighter than predicted \citep[e.g.,][]{blaesetal2001}. The X-ray size has been shown by microlensing to be extremely compact \citep{chartasetal2009,daietal2010}, disfavoring spectral models invoking global Comptonizing atmospheres to explain the UV emission. Quasar optical variability is also difficult to explain in the context of the thin disk model. Long term monitoring of AGN has found almost simultaneous variability across optical wavelengths, with lags of less than 1-2 days \citep{cutrietal1985,claveletal1991,koristaetal1995,wandersetal1997,collieretal1998}. Comparing these lags to the radii dominating the thin disk emission at these wavelengths gives a traveling speed of 0.1c for the variability mechanism \citep{kroliketal1991,courvoisierclavel1991}. This would force the variability in AGN to be communicated at the local sound or Alfv\'{e}n speed rather than on the infall timescale associated with disk instabilities. \citet{wambsganssetal1990} and \citet{rauchblandford1991} first measured the size of an accretion disk from microlensing, and found that accretion disks were smaller than expected from the thin disk model for the observed $L_\nu$. In recent years, the opposite trend has emerged \citep{pooleyetal2007,daietal2010,morganetal2010,blackburneetal2010}. Sizes are now robustly found to be a factor of $\sim4$ larger on average than expected from the thin disk model at IR/optical/UV wavelengths. The size vs. wavelength relation predicted by the thin disk, $r \propto \lambda^{4/3}$, is within the large range allowed by microlensing observations \citep{eigenbrodetal2008}. At the same time, \citet{kellyetal2009}, \citet{kozlowskietal2010} and \citet{macleod2010} have studied large samples of quasar light curves. They find that optical quasar variability is well described by a damped random walk which returns to a mean value on a typical timescale of 200 days with variability amplitudes of $\simeq 10-20\%$. It is unlikely that the observed quasar variability is caused by a coherently varying accretion disk\footnote{The implications of coherent variations for microlensing measurements are discussed by \citet{blackburnekochanek2010}.}, but rather is probably the added effect of many smaller, independently varying regions. Many such models have been proposed to explain quasar variability \citep[e.g.][]{lyubarskii1997}. In this Letter, we demonstrate that for the observed variability characteristics, such an inhomogeneous disk can simultaneously explain multiple discrepancies between AGN observations and accretion disk theory. Inhomogeneous disks can be large enough to explain the microlensing observations, while their temperature fluctuations on small spatial scales naturally explain the observed simultaneous variability across optical wavelengths. Temperatures exceeding the local thin disk value lead to broader spectra extending into the UV, consistent with quasar spectra without invoking a Compton scattering medium. \begin{figure} \epsscale{1.2} \plotone{f2.eps} \caption{\label{rarrsdmn}Median size increase vs. relative variability for various zone sizes for the damped random walk model. Each sequence has values of $\sigma_T$ from $0.1-0.8$. The allowed region from microlensing and variability studies is shaded. The lines show analytic fits to the variability from the damped random walk model (Eq. \ref{varfit}), and to the fractional increase in half-light radius from the log-normal model (Eq. \ref{sizefit}). For $n \gtrsim 100$ and $\sigma_T \simeq 0.4$, damped random walk disks can be large enough to explain microlensing while matching the observed variability.} \end{figure}

Standard thin disk accretion has formed the basis for understanding X-ray binaries and AGN for nearly 40 years. However, it has always been difficult to extend the model to explain optical quasar variability. In recent years, microlensing observations of multiply imaged quasars have provided a probe of the disk structure, finding that quasar microlensing sizes are robustly larger than the flux sizes predicted from thin disk theory. If the average temperature structure remains identical to that in thin disk theory but is highly inhomogeneous, accretion disks can be large enough to explain the sizes found by microlensing while matching the observed level of optical variability. The level of inhomogeneity ($\sigma_T=0.35-0.50$ with $n=10^{2-3}$) required to explain the discrepancy in microlensing sizes is in excellent agreement with that necessary to fit observed quasar spectra. Such inhomogeneous structure produces short timescale variations in microlensing light curves that should be larger than $\simeq 1\%$. The amplitude of the temperature fluctuations can be further constrained from measuring the size of these deviations. The range of temperatures in small regions of the inhomogeneous disk explains the simultaneous variability observed across optical wavelengths. We have demonstrated this idea with an unphysical toy model. Proper modeling of an inhomogeneous disk will require global MHD simulations of radiation-dominated accretion disks. It is unclear whether the MRI alone is sufficient to produce the required temperature fluctuations, or whether additional disk instabilities or other variability mechanisms are also important.

1012.3169

1012

1012.1393_arXiv.txt

We phenomenologically developed a propagation model of high energy galactic cosmic rays. We derived the analytical solutions by adopting the semi-empirical diffusion equation, proposed by Berezinskii {\it et al.}(1990) and the diffusion tensor proposed by Ptuskin {\it et al.}(1993). This model takes into account both the symmetric diffusion and the antisymmetric diffusion due to the particle Hall drift. Our solutions are an extension of the model developed by Ptuskin {\it et al.} (1993) to a two-dimensional two-layer (galactic disk and halo) model, and they coincide completely with the solution derived by Berezinskii {\it et al.} (1990) in the absence of antisymmetric diffusion due to Hall drift. We showed that this relatively simple toy model can be used to explain the variation in the exponent of the cosmic ray energy spectrum, $\gamma$, around the knee $E \approx 10^{15}$ eV.

} Cosmic ray propagation is one of the most important and interesting subjects in astrophysics and high energy particle physics. It is believed that the observed cosmic ray data such as the cosmic ray energy spectrum includes information about the space through which the cosmic rays pass. In fact, it is possible to evaluate the thickness of matter, and it is also thought that details about the galactic magnetic field can be extracted from the observed data because the cosmic rays experience frequent collision and scattering with both the interstellar gas and the galactic magnetic field during their propagation in the galaxy. Thus, a reliable cosmic ray propagation model may enable us to obtain further knowledge about the galactic structure. Thus far, several cosmic ray propagation models have been proposed and discussed (see \cite{cesarsky,ptuskin2001} and references therein). Further, some researchers have successfully derived analytical solutions using the diffusion equation \cite{berezinskii,pacheco,guet,ptuskin}. The spectrum of cosmic rays shows one of the most distinctive features, known as ``knee'', around energy $E \approx 10^{15}$ eV at which the exponent of the energy spectrum, $\gamma$, changes from 2.6 -- 2.7 for $10^{10} \le E \le 10^{15}$ eV to 3 -- 3.1 for $10^{15} \le E \le 10^{18}$ eV. Currently, it is not clear why the exponent changes around $E \approx 10^{15}$ eV. Thus far, several models have been proposed to explain this spectral property: a shock wave acceleration model based on the acceleration of cosmic ray particles by the shock wave front \cite{berezhko,stanev,kobayakawa,sveshnikova,erlykin,volk,plaga}, a diffusive propagation model based on the leakage and diffusive propagation of cosmic rays in the galaxy \cite{swordy,lagutin,ptuskin,ogio,roulet,candia1,candia2,candia4}. an interaction model based on the interaction of cosmic rays with the background particles in the galaxy \cite{tkaczyk,karakula,dova,candia3}, a reaction model based on the reaction of cosmic rays with the atmosphere of Earth \cite{kazanas1,kazanas2}, etc. Among these models, the shock wave acceleration model seems to be widely accepted as the explanation for the knee. However, the diffusive propagation model can provide the exposition for the first knee and the second knee and for the observed compositions and anisotropies \cite{ptuskin,candia1,candia2,candia4}. The diffusive propagation model is characterized by introducing the particle Hall drift effect; thus, it is a theoretically simple model. In this study, we derived solutions for the diffusive propagation model of cosmic rays and confirmed the validity of this model. We adopted the semi-empirical diffusion equation introduced in \cite{berezinskii} and the diffusion tensor described in \cite{ptuskin}. Then, we extended the propagation model in \cite{ptuskin} to a two-dimensional two-layer (comprising the galactic disk and the halo) model in cylindrical coordinates. This paper is organized as follows: In Section \ref{solution}, we derive the analytical solutions for the diffusion equation of cosmic rays. In Section \ref{psols}, we qualitatively show that our model can explain the spectral feature of the observed cosmic rays, namely, the exponential variation around the knee. In Section \ref{conclusion}, we present the conclusions of our study.

} We phenomenologically proposed a propagation model of galactic cosmic rays based on the semi-empirical diffusion equation developed by Berezinskii {\it et al.} (1990) and the diffusion tensor introduced by Ptuskin {\it et al.} (1993). This model takes into account both the symmetric diffusion and the antisymmetric diffusion due to the particle Hall drift. The derived solutions are an extension of the model developed by Ptuskin {\it et al.} (1993) to a two-dimensional two-layer (galactic disk and halo) model, and they coincide completely with the solutions derived by Berezinskii {\it et al.} (1990) in the absence of antisymmetric diffusion due to particle drift. We shown that this relatively simple model can be used to explain the variation in the exponent of the cosmic ray energy spectrum, $\gamma$, around the knee $E \approx 10^{15}$ eV. In this paper, we showed that although the diffusive cosmic ray propagation model can be used to explain the observed cosmic ray spectrum, especially the exponential variation around the knee, our model is actually a more simple toy model based on assumptions such as the cylindrical structure of the galaxy and simplification of magnetic field. To further test the validity of the diffusion model, we must conduct numerical simulations under more realistic situations. This is a difficult task; nonetheless, it may help us to gain a deeper understanding of astroparticle physics and the galactic structure through which the cosmic rays pass.

1012.1393

1012

1012.5673_arXiv.txt

The historical supernova remnant (SNR) Tycho SN 1572 originates from the explosion of a normal Type Ia supernova which is believed to have originated from a carbon-oxygen white dwarf in a binary system. We analyze the 21cm continuum, HI and $^{12}$CO-line data from the Canadian Galactic Plane Survey in the direction of SN 1572 and surrounding region. We construct HI absorption spectra to SN 1572 and three nearby compact sources. We conclude that SN 1572 has no molecular cloud interaction, which argues against previous claims that a molecular cloud is interacting with the SNR. This new result does not support a recent claim that dust, newly detected by AKARI, originates from such a SNR-cloud interaction. We suggest that the SNR has a kinematic distance of 2.5 - 3.0 kpc based on a nonlinear rotational curve model. Very-high-energy $\gamma$-ray emission from the remnant has been detected by the VERITAS telescope, so our result shows that its origin should not be an SNR-cloud interaction. Both radio and X-ray observations support that SN 1572 is an isolated Type Ia SNR.

Recent observations have revealed that the historical Supernova Remnant (SNR) Tycho SN 1572 belongs to the class of Type Ia SN by detecting its optical spectrum near maximum brightness from the scattered-light echo \citep{Kra08}. A red subgiant has been suggested to be the possible surviving companion of the supernova in a close binary system \citep{Rui04}. However, the evolutionary path of the progenitor is still not understood, and this association has been questioned \citep{Fuh05, Iha07}. SN 1572 is a natural candidate for high energy observations. Non-thermal X-ray emission and thin filamentary structures in the remnant are believed to be associated with high-energy electron acceleration \citep{Bam05, Kat10}. \citet{War05} studied the shock dynamics using Chandra observations and suggested there is efficient hadronic particle acceleration in the remnant. Weak TeV emission coming from the direction of SN 1572 is newly detected by the VERITAS telescope \citep{Acc10}, although confirmation by other instruments (e.g. MAGIC) is required. Young, massive core-collapse SNe are usually not far away from their parent molecular clouds since their progenitors evolve very quickly (a few million years). Therefore, it is expected that many Galactic Type II/Ibc SNRs are associated with large molecular clouds. Recently, \citet{Jia10} cataloged more than 60 possible SNR-cloud interaction systems. As the only known Type Ia remnant in this catalogue, SN 1572 has been proposed to be interacting with dense ambient (atomic/molecular) clouds toward its Northeast (NE) \citep{Rey99, Lee04, Cai09}. Cold dust overlapping the eastern part of SN 1572 has been detected \citet{Ish10} and taken as evidence of a possible interaction between SN 1572 and a molecular cloud. Extended TeV emission detected in several SNRs has been suggested to originate from the interaction between the SNR shock and an adjacent CO cloud \citep{Eno02, Aha04, Alb07, Tia08, Cas10, Tav10}. Is TeV emission from SN 1572 associated with such a SNR-cloud interaction? In this letter, we take advantage of the 21cm continuum, HI and $^{12}$CO-line data from the Canada Galactic Plane Survey (CGPS) in the direction of SN 1572. We study if there exists such an interaction responsible for the TeV emission. Details on the CGPS are given in \citet{Tay03}, the analysis methods are described in our previous papers \citep{Tia07, Tia10}.

\subsection{The CO cloud at -64 km s$^{-1}$ is behind SN 1572} We have found that the HI gas at velocities of -47 to -53 km s$^{-1}$ is in front of SN 1572. Both the CO at -64 km s$^{-1}$ and the HI at -60 km s$^{-1}$ do not produce associated HI absorption, so both are likely behind SN 1572. Is it possible that the CO cloud could still be in front of SN 1572 but the amount of cold HI gas in the molecular cloud is too small to produce detectable HI absorption against SN 1572? Fig. 2a shows that the CO at -64 km s$^{-1}$ has associated HI emission (T$_{B}$$\sim$ 75 K). Technically, a minimum optical depth of $\tau$ $\sim$ 0.1 may be detected. This requires that HI atomic gas against background continuum source has at least column density of N$_{HI}$=2.3$\times$10$^{19}$ cm$^{-2}$ (N$_{HI}$=1.823$\times$10$^{18}$ $T_{s}$$\tau$$\Delta{v}$; \citet*{Dic90}), given $T_{s}$$\sim$ 25 K (\citep{Sch95} and the FWHM $\Delta{v}$ of the typical HI absorption line $\sim$ 5 km/s. Taking the theoretical mean HI/H$_{2}$ ratio of 0.2 \citep*{Gol05, And09}, this respective molecular cloud has a H$_2$ column density of $\sim$ 1.2$\times 10^{20}$ cm$^{-2}$. The $^{12}$CO spectrum from box 1 (Fig. 2a) gives W$_{^{12}CO}$ $\sim$ 7 K km s$^{-1}$ for the cloud component. Taking the $^{12}$CO to H$_{2}$ conversion factor of $X$=3$\times$10$^{20}$ [cm$^{-2}$/(K km s$^{-1}$)], this gives a H$_{2}$ column density of $\sim$ 2.1$\times$10$^{21}$ cm$^{-2}$ for the respective cloud. This is enough to produce measurable HI absorption at 18 $\sigma$ level if the cloud component is in front of SN 1572. So the CO molecular cloud at -64 km s$^{-1}$ is behind SN 1572. \subsection{Does there exist SN 1572-CO cloud interaction?} Previous multi-band observations of SN 1572 \citep{Str82, Gha00, Hug00, Dou01, Hwa02, Lee07, Yan09} have revealed limb-brightened radio and X-ray shell, H$\alpha$ filament along the NE boundary of the shell, the mid-IR emission from SN 1572, suggesting that SN 1572 is surrounded by an inhomogeneous environment. A possible SNR-cloud interaction along SN 1572's NE boundary is able to trigger some observed phenomena, e.g. the NE bright part in radio and X-ray images, the decelerated expansion of the NE rim in radio, optical and X-rays, the dust emission at the NE boundary, etc. We have concluded that the CO cloud at -64 km s$^{-1}$ is behind SN 1572, but is it possible that the CO is adjacent and interacting with SN 1572? This cloud has a H$_{2}$ column density of 2.1 $\times$10$^{21}$ cm$^{-2}$, so its density is $\sim$ 200 cm$^{-3}$, assuming half of the SNR is surrounded by the cloud (R$_{SN 1572}\sim$ 4 arcmin) and taking a distance of 3 kpc (see section 3.5 for detail). This is inconsistent with recent Chandra X-ray observations of the cloud surrounding the remnant. \citet{Cas07} and \citet{Kat10} analyzed X-ray spectra from the thin rim between the blast wave and contact discontinuity, and found that there is little thermal emission from the preshock ambient medium. This requires that the ambient medium density in the vicinity of SN 1572 is less than 0.2 cm$^{-3}$. This is three orders of magnitude lower than the density of the CO cloud. Therefore the CO cloud is not adjacent to SN 1572. We conclude that there is no physical association between SN 1572 and the $^{12}$CO cloud at -64 km s$^{-1}$, although they are overlapping along the light-of-sight. \citet{Ish10} detected cold dust IR emission outside the NE and Northwest (NW) boundaries of SN 1572's shell, and suggested that the NE dust emission comes from a possible molecular cloud interacting with the shock front. The origin of the NW dust emission is rather unclear because of the absence of any interstellar cloud nearby. However, our study reveals that the origin of the NE dust emission is also unclear, probably from the molecular cloud at -64 km s$^{-1}$ but unrelated to SN 1572. We notice that weak TeV emission from SN 1572 is also detected \citep{Acc10}. Although TeV emission from several SNRs has been suggested to originate from interaction between the SNR shock and an adjacent cloud, our result reveals it is not this case for SN 1572 at least. \subsection{Does there exist SN 1572-HI cloud interaction?} \citet{Rey99} studied 21 cm spectra in the velocity range of -41 to -106 km s$^{-1}$ towards SN 1572 using Very Large Array (VLA) archive data and single-dish HI observations. They detected HI absorption from -46.4 to -56.8 km s$^{-1}$ towards SN 1572, and found an extended HI absorption along the eastern side of the shell between the velocities of -47.7 and -50.3 km s$^{-1}$. They also found a small high-density HI clump (160-325 cm$^{-3}$) observed as an absorption feature at -51.5 km s$^{-1}$ towards the eastern part of the shell. Our study reproduces some of these results: The extended HI structure and the small HI clump seen in our Fig. 4 (at velocities of -47.63 and -52.58 km s$^{-1}$) are found. By examining the HI channel maps from -50.3 to -60.6 km s$^{-1}$ shown in their Fig. 2, we see a clear deficit of HI brightness between -52.9 and -56.8 km s$^{-1}$ surrounding SN 1572. This is inconsistent with being a HI absorption feature because it does not correspond with bright continuum emission. It could be either HI self absorption (HISA) or artifacts. Our HI absorption spectra clearly reveal reliable HI absorption features in the velocity range of -47 to -53 km s$^{-1}$ but not beyond -53 km s$^{-1}$. Our methods to build HI absorption spectra have reduced false absorption features as much as possible. We do not find any reliable absorption towards SN 1572 in our HI channel maps in the range of -54 to -66 km s$^{-1}$. So any HI absorption feature beyond -53 km s$^{-1}$ is likely not real. This extended HI cloud along the eastern side of the shell is in front of SN 1572, but is it possible to be adjacent to SN 1572? The HI cloud's density of $\sim$10 cm$^{-3}$ ($\tau$$\sim$0.8, $\Delta$$v$$\sim$3 km s$^{-1}$ from Fig. 2a and 3, T$_{s}$=25K from \citet{Sch95}) is much higher that the ambient medium density of 0.2 cm$^{-3}$ from the X-ray measurements, so the HI cloud is not adjacent to SN 1572. \citet{Rey99} suggested an interaction between SN 1572 and the small NE HI clump because of two factors: the HI clump is near the site of the lowest expansion velocity along the whole shell of SN 1572 (which can roughly explain the slowing of expansion rate of the eastern rim), and a detected H$\alpha$ knot lies near the HI clump. However, \citet{Lee07}, using the Subaru Telescope, obtained a systemic velocity of -30.3$\pm$0.2 to the H$\alpha$ knot G. So we need further new independent observations to distinguish if the knot has relation with the HI clump or not. In summary, there is no direct evidence that the extended HI cloud along the eastern part of SN 1572 is physically associated with SN 1572. \subsection{Tycho SN 1572, a naked Ia SNR} Massive stars have a lifetime of about 10$^{6}$ yr, so Type II/Ibc SNe are expected to take place in a dense, star-forming region where their progenitors are formed. This is different for Type Ia SNe because of the longer time needed for the system to evolve ($\sim$10$^{8}$ yr). SN 1572's progenitor has wandered $\sim$1 kpc far away from the star-forming site, given an average birth velocity of 10 km s$^{-1}$, therefore is far outside of its parent molecular cloud (generally the size of individual giant molecular cloud is $\sim$ 100 parsec). Although SN 1572 might encounter other clouds as it wanders through the interstellar medium, X-ray observations of 1572 show low ISM density. So we believe SN 1572 is likely a naked SNR. \subsection{Distance to Tycho SN 1572} The distance to SN 1572 has previously been suggested between 2 to 5 kpc by radio, optical, X-ray and $\gamma$-ray observations (Fig. 6 of \citet{Hay10} shows a summary). As the Perseus arm is influenced by the spiral shock (leading to a velocity reversal; \citet{Rob72}), it is challenging to estimate kinematic distances to objects in the Perseus arm of the outer Galaxy. This velocity reversal causes a distance ambiguity for gas and objects with radial velocities of -40 km s$^{-1}$ to -55 km s$^{-1}$ in the line-of-sight to SN 1572, for $v \le$ -55 km s$^{-1}$ the radial velocity decreases monotonically with distance (see Fig. 2 of \citet{Alb86}). Previously, HI absorption observations have been made towards SN 1572. \citet{Alb86} made aperture synthesis observations of HI using the Cambridge Half-Mile Telescope, and suggested a distance in the range 1.7 - 3.7 kpc. \citet{Sch95} used the VLA to study HI absorption towards SN 1572 and nearby compact sources including G120.56+1.21 and G119.74+2.4, and estimated a distance of 4.6$\pm$0.5 kpc. The distance difference between \citet{Alb86} and \citet{Sch95} is caused by how to deal with a possible weak absorption feature at -60 km s$^{-1}$. \citet{Alb86} thought the absorption could be either from HI in a turbulent state around a filament or could be spurious caused by small-scale variation in HI emission. \citet{Sch95} believed it is real and SN 1572 is farther than the region of distance ambiguity. We obtain HI absorption features with higher quality than previous studies, and find clearly that the highest absorption velocity is -53 km s$^{-1}$ towards SN 1572. Therefore we exclude the large distance of 4.6 kpc. Due to absence of HI absorption between -40 and -45 km s$^{-1}$ towards SN 1572, \citet{Alb86} further proposed that the HI emission between -40 and -45 km s$^{-1}$ could be behind SN 1572, and that SN 1572 most likely is located at the near distance of the major absorption feature at -50 km s$^{-1}$. Anyway, they kept an option that SN 1572 has small probability to be at the far side distance of 3.7 kpc of the same absorption feature. Figure 2 show that the HI gas at -41 to -46 km s$^{-1}$ has no associated obvious HI absorption towards SN 1572 and definitely produces associated HI absorption towards G120.56+1.21. We notice that the HI channel maps in this velocity range show more HI gas with brightness temperature above 20 K surrounding G120.56+1.21 than SN 1572. Could absence of HI absorption in the velocity range towards SN 1572 be due to insufficient HI gas to produce measurable optical depth in front of SN 1572? However, the HI gas has column density of $\sim$ 3$\times$10$^{20}$ cm$^{-2}$ (N$_{HI}$=1.832$\times$10$^{18}$$T_{B}$$\Delta{v}$, $T_{B}$=35 K, see Fig. 2a) which is enough to produce detectable HI absorption, because $\tau$$\ge$ 0.1 requires a minimum HI column density of $\sim$ 7$\times$10$^{19}$ cm$^{-2}$ (here T$_{s}$=75 K, which is an average value in the low velocity range from \citet{Sch95}. They also noticed that it goes below 25 K near -50 km s$^{-1}$). So this gives an upper limit distance for SN 1572 which is in front of the HI at -41 to -46 km s$^{-1}$. In addition, \citet{Sch95} detected HISA at -49 km s$^{-1}$. HISA is generally produced by foreground cold HI in front of background warm HI at same velocity. Because of the velocity reversal in the Perseus Arm, HISA features are widely observed in the CGPS \citep{Gib05, Tia10}. The HISA at -49 km s$^{-1}$ likely originates from the same cause, i.e. cold HI at near side of the velocity reversal absorbs emission from warm HI at the far side at same velocity of -49 km s$^{-1}$. Because the - 49 km s$^{-1}$ HI is in front of SN 1572 (Figs. 2 and 3), SN 1572 must be between the HI at at near side with -48 km s$^{-1}$ and the HI at far side with -41 to -46 km s$^{-1}$. In other words, the distance to SN 1572 is between 2.5 to 3.0 kpc. We use the \citet{Fos06}'s model which is similar with the Robert's model but puts the spiral shock front at 2.5 kpc in the direction to SN 1572 (also see Fig. 14 of \citet{Sch95}) A kinematic distance of 2.5 to 3.0 kpc is roughly consistent with new estimates from other independent methods. \citet{Vol08} suggested that SN 1572's distance is greater than 3.3 kpc by modeling the existing $\gamma$-ray measurements from SN 1572. \citet{Kra08} estimated distance of 3.8$^{1.5}_{-0.9}$ kpc using classic brightness-distance relation and accounting for interstellar foreground extinction. \citet{Hay10} made new $Suzaku$ observations of SN 1572 and estimated an average spherical expansion velocity of $\sim$4700 km s$^{-1}$. They gave a direct distance estimate of 4$\pm$1 kpc by combining the observed ejecta velocities with the ejecta proper-motion measurement by $Chandra$.

1012.5673

1012

1012.5359_arXiv.txt

{ We are now exploring the inner region of Type 1 active galactic nuclei (AGNs) with the Keck interferometer in the near-infrared. Adding to the four targets previously studied, we report measurements of the K-band (2.2~$\mu$m) visibilities for four more targets, namely AKN120, IC4329A, Mrk6, and the radio-loud QSO 3C273 at $z$=0.158. The observed visibilities are quite high for all the targets, which we interpret as an indication of the partial resolution of the dust sublimation region. The effective ring radii derived from the observed visibilities scale approximately with $L^{1/2}$, where $L$ is the AGN luminosity. Comparing the radii with those from independent optical-infrared reverberation measurements, these data support our previous claim that the interferometric ring radius is either roughly equal to or slightly larger than the reverberation radius. We interpret the ratio of these two radii for a given $L$ as an approximate probe of the radial distribution of the inner accreting material. We show tentative evidence that this inner radial structure might be closely related to the radio-loudness of the central engine. Finally, we re-observed the brightest Seyfert 1 galaxy NGC4151. Its marginally higher visibility at a shorter projected baseline, compared to our previous measurements obtained one year before, further supports the partial resolution of the inner structure. We did not detect any significant change in the implied emission size when the K-band flux was brightened by a factor of 1.5 over a time interval of one year.}

1012.5359

1012

1012.2431.txt

% insert abstract here We establish a new self-consistent model in order to explain from a unified viewpoint two key features of the cosmological evolution: the inflation in the early Universe and the late-time accelerated expansion. The key element of this new model is the Archimedean-type coupling of the dark matter with dark energy, which form the so-called cosmic dark fluid. We suppose that dark matter particles immersed into the dark energy reservoir are affected by the force proportional to the four-gradient of the dark energy pressure. The Archimedean-type coupling is shown to play a role of effective energy-momentum redistributor between the dark matter and the dark energy components of the dark fluid, thus providing the Universe evolution to be a quasiperiodic and/or multistage process. In the first part of the work we discuss a theoretical base and new exact solutions of the model master equations. Special attention is focused on the exact solutions, for which the scale factor is presented by the anti-Gaussian function: these solutions describe the late-time acceleration and are characterized by a nonsingular behavior in the early Universe. The second part contains qualitative and numerical analysis of the master equations; we focus there on the solutions describing a multi-inflationary Universe.

The concepts of dark energy (DE) and dark matter (DM) (see, e.g., \cite{DE1,DE2,DE3} and \cite{DM1,DM2,DM3} for review and references) are the basic elements of modern cosmology and astrophysics. These elements were introduced into the scientific lexicon in two different ways: dark energy is considered to be a reason for the late-time accelerated expansion of the Universe \cite{A1,A2}, while dark matter is usually associated with the explanation of the flat velocity curves of the spiral galaxies rotation \cite{Sp1,Sp2}. Nevertheless, there exists a tendency to consider dark energy and dark matter as two manifestations of one unified dark fluid (see, e.g., \cite{DF1,Oreview,DF2,DF3,DF4,DF5}). The models of interaction between two constituents of the dark fluid, as well as the models of interactions of dark energy and/or dark matter with the standard (baryon) matter, are the subjects of wide discussion. The total contribution of dark energy and dark matter into the energy balance of the Universe is estimated to be about $95 \%$. Thus, the coupling between these two constituents of the dark fluid seems to be the most important element in the list of cosmic medium interactions, and the dark fluid can be considered as a thermodynamic reservoir for baryon matter. The most developed model of the coupling between DM and DE constituents of the dark fluid is based on the two-fluid representation of the cosmic medium (see, e.g., \cite{M1,M2,M3,M4,M5,M6}). In this approach the interaction terms $\pm Q$ appear in the right-hand sides of separate balance equations for the DE and DM with opposite signs and disappear in a sum, when one deals with the total balance equation. The modeling of the coupling term $Q$ has in most cases a phenomenological character and is based on the ansatz that $Q$ is a function (e.g., linear or power-law) of the energy densities of the DE and DM, of the Hubble function, $H$, etc. We formulate the theory of interaction between DE and DM using the relativistic hydrodynamics for dark energy and relativistic kinetics for dark matter. We suggest that the DE acts on the DM particles by means of some effective force, and the corresponding model force four-vector is introduced into the kinetic equation. The backreaction of the DM on the DE is described using a force-type term in the hydrodynamic equations. The total system (DM plus DE) is considered to be conservative. The concept of the Archimedean-type force can be naturally generalized for the description of the DE action on the baryon matter; however, now we restrict ourselves by the model of the DE and DM interaction. We discussed the structure and properties of various effective forces, which appeared in the cosmological contexts, in the papers \cite{BZ1,BZ2,BZ3,BZ4,BZ5,BZ6}, the relativistic generalizations of the Stokes force, Langevin force, antifriction and tidal forces being investigated in detail. Concerning the force acting on the DM particles from the DE we would like to introduce the so-called Archimedean-type force. This choice can be motivated as follows. {\it First} of all, this force is a relativistic generalization of the classical Archimedean force, proportional to the three-gradient of the Pascal pressure; thus, the model under discussion is based on the well-known and well examined scheme of interaction. {\it Second}, the DE pressure is assumed to be of the same order as the DE energy density (73\% of the Universe energy density), thus, the Archimedean effect on the dark matter could be significant. {\it Third}, assuming that the DE pressure can be negative, we obtain that such a force can be attractive in contrast to expulsive classical Archimedean force. Admitting that the DE pressure changes the sign in the course of the Universe evolution, one can describe a multistage (or even quasiperiodic) character of the cosmological expansion, for which epochs of deceleration are changed by epochs of acceleration and vice versa. (The interest in models of this type was renewed by the paper \cite{Penrose}). We show that, in principle, the Archimedean-type force can effectively redistribute $95\%$ of the Universe's energy between the DE and DM constituents, thus guiding the time evolution of the cosmic medium as a whole. We divided the work into two parts: the first one contains pure analytical results and exact solutions of the model; in the second part we focus on the numerical and qualitative analysis of the model. The first part of the work is organized as follows. In Sec.II we derive the master equations of the model with an Archimedean-type force. In particular, in Sec. II.B based on the kinetic approach we introduce the Archimedean-type force, obtain basic integrals of motion, construct the distribution functions and calculate their macroscopic moments as functions of the DE pressure. In Sec. II.C we formulate the balance equation for the dark fluid. In Sec. II.D we discuss the extended (inhomogeneous) equation of state for the dark energy and introduce the key equation for the DE pressure evolution. Sec. III contains discussions about two classes of exact solutions. In Sec. III.A we consider a special (constant) exact solution for the case when the guiding parameter of the model $\sigma$ is not equal to its critical value, i.e. $\sigma \neq -1$, and we analyze the problem of asymptotic stability of this solution. In Sec. III.B we focus on the special case $\sigma = -1$ and obtain exact solutions of the anti-Gaussian type for the following submodels: (i) the massless DM, (ii) the cold dark matter, and (iii) the submodel with the DE domination. In Sec. IV we discuss obtained analytical results.

The study of the cosmological model, into which the Archimedean-type interaction between dark energy and dark matter is introduced, shows that the roles of DM and DE in the energy balance of the Universe can be revised. According to the obtained formula (\ref{E(x)}), the contribution of the DM particles, $E_{({\rm a})}(a(t))$, into the total energy density depends on the state of DE pressure $\Pi(a(t))$ through the functions $F_{({\rm a})}(x)$ (see (\ref{Fdefin})). In the models without Archimedean-type force the energy of the DM particle decreases effectively because of the factor $a^{-2}(t)$; in other words, all the particles, both nonrelativistic and ultrarelativistic at the initial moment $t_0$, inevitably become (effectively) nonrelativistic in the process of Universe expansion. When the Archimedean-type force acts on the DM particles, the particle energy (\ref{energy}) becomes much more complicated function of cosmological time due to the exponential dependence on the DE pressure. This means, in particular, that, nonrelativistic particles can become (effectively) ultrarelativistic due to the Archimedean-type force action, thus the corresponding contribution of cold dark matter into the total energy can be reestimated taking into account the sign, the value of the DE pressure at this moment, as well as the rate of its variation with time. In contrast, the ultrarelativistic DM particles can become (effectively) nonrelativistic, when the corresponding exponential factor in (\ref{energy}) is small. Now the contribution of cold dark matter into the total energy is estimated to be about $23 \%$. The question arises: does this estimate include a total rest energy of the massive DM particles only, or the energy of the Archimedean-type interaction as well? We need such a clarification, for instance, in order to estimate the density numbers of the DM particles of different sorts; in particular, the information about the number density of DM axions in the terrestrial laboratories is very important for planning experiments with axion electrodynamics (see, e.g., \cite{WTNi}). \vspace{3mm} The cosmological model with Archimedean-type force describes a self-regulating Universe. This means that in the process of expansion of the Universe the total (conserved as a whole) energy can be redistributed between dark energy and dark matter constituents according to the challenge of the corresponding epoch. The energy pendulum stimulated by the Archimedean-type force can work by the following scheme: let us imagine that at some moment the DE pressure $\Pi$ is negative and is varying rather quickly; then according to the formulas (\ref{key2}), (\ref{key3}), (\ref{nrP00key}) the DM particle reaction, provoked by the Archimedean-type force, will be strong. The corresponding intensive source appears in the right-hand side of the key equation (\ref{key1}), thus decreasing the rate of $\Pi$ evolution. From the theoretical point of view, it is not yet clear, first, for which set of guiding parameters such a specific regime does exist; second, when such a regime can be characterized as (quasi)oscillations; and third, how the number of epochs of the Universe expansion does depend on the set of guiding parameters of the model. We started to study these questions qualitatively and numerically in the second part of our work and presented examples of quasi-periodic, multi-inflationary and multistage evolutionary schemes in the framework of the model based on the Archimedean-type interaction between dark energy and dark matter. \vspace{3mm} In the first part of the work we focused on exact solutions of this new model. The first exact solution is the constant one, $\Pi(x) {=} {-} \rho(x) {=} {-}\frac{\rho_0}{1{+}\sigma}$, and relates to the case $\sigma \neq {-}1$. Since the DE pressure for this exact solution is constant, the Archimedean-type force becomes hidden, and we obtain the standard cosmology with $\Lambda$ term. This solution is asymptotically stable, when the guiding parameters of the model, $\xi$ and $\sigma$, satisfy the inequalities $0<\xi<\frac{1}{3}$ and $\sigma > {-}3\xi$, or $\xi \geq \frac{1}{3}$ and $\sigma > {-}1$. When $\sigma < {-}1$, the solution is asymptotically unstable for arbitrary parameter $\xi$. \vspace{3mm} Exact solutions of the second class, with $\sigma=-1$, are much more sophisticated, since the DE pressure is described by the logarithmic function of the ratio $a(t)/a(t_0)$. The corresponding scale factor $a(t)$ is given by the anti-Gaussian function (\ref{toy5}), it has no singular points in the early Universe, and describes late-time accelerated expansion with the acceleration parameter ${-}q>1$; this acceleration parameter is bigger than for the de Sitter model. Exact solutions of the anti-Gaussian type happen to be typical for different physical situations: for massless and massive nonrelativistic DM, for the case with DE domination, etc. This solution is unstable. \vspace{3mm} There are two specific values of the guiding parameters of the model: $\sigma {=} {-}1$ and $\xi {=} \frac{1}{3}$. The first one, $\sigma {=} {-}1$, can be associated with the so-called phantom divider $w(t) {=} {-}1 {=} \frac{1}{\sigma}$. The second value, $\xi {=} \frac{1}{3}$, can be denoted as a resonance value of the relaxation time parameter. Indeed, according to the formula (\ref{simplest0}) the function $\frac{\xi}{H(t)}$ plays a role of the relaxation time for the DE pressure $\Pi(t)$, say, $\tau_{\Pi}$. When $\xi {=} \frac{1}{3}$, one obtains that $\tau_{\Pi} {=} \frac{1}{3H(t)} {=} \frac{1}{\Theta(t)}$, where $\Theta(t) {=} 3 H(t) {=} \nabla_k U^k$ is the expansion parameter. Thus, the characteristic time of expansion $\frac{1}{\Theta(t)}$ coincides with the relaxation time parameter for the DE pressure, introducing some specific resonance condition. When $\sigma {=} {-}1$ and $\xi {=} \frac{1}{3}$ simultaneously, there exists superexponential solution of the model, described by the scale factor (\ref{Sc61}). \vspace{3mm} A physical origin of the Archimedean-type force is not yet clear; nevertheless, this force seems to be very interesting from the viewpoints of a new model of interaction and a new scheme of redistribution of the cosmic energy between the interacting DE and DM constituents. The presented model is self-consistent, simple from the point of view of analysis and very promising. We keep in mind the story of the Chaplygin gas model \cite{Chap}: starting from a classical analogy a new evolutionary model has been elaborated and applied to cosmology, although physical meaning of the Chaplygin scheme of interaction is under discussion till now. \vspace{5mm}

1012.2431

1012

1012.2862_arXiv.txt

We searched for radio pulsars in 25 of the non-variable, unassociated sources in the \fermi\ LAT Bright Source List with the Green Bank Telescope at 820\,MHz. We report the discovery of three radio and \gray\ millisecond pulsars (MSPs) from a high Galactic latitude subset of these sources. All of the pulsars are in binary systems, which would have made them virtually impossible to detect in blind \gray\ pulsation searches. They seem to be relatively normal, nearby ($\le$2\,kpc) millisecond pulsars. These observations, in combination with the \fermi\ detection of \grays\ from other known radio MSPs, imply that most, if not all, radio MSPs are efficient \gray\ producers. The \gray\ spectra of the pulsars are power-law in nature with exponential cutoffs at a few GeV, as has been found with most other pulsars. The MSPs have all been detected as X-ray point sources. Their soft X-ray luminosities of $\sim$10$^{30-31}$\,erg\,s$^{-1}$ are typical of the rare radio MSPs seen in X-rays.

Before the launch of the \fermi\ Gamma-ray Space Telescope, the only pulsars with definitive detections in \grays\ (using EGRET on {\it CGRO}) were young and very energetic ($\dot E>10^{36}$\,erg\,s$^{-1}$) or nearby older systems ($\dot E>10^{34}$\,erg\,s$^{-1}$) \citep{thompson04}. A possible detection of pulsed \grays\ from the energetic MSP J0218$+$4232 \citep{egret0218}, sparked interest in modeling MSP \gray\ emission \citep[e.g.][]{zc03,hum05}, and encouraged one group \citep{sgh07} to predict that many new MSPs might be detected in \grays\ or discovered in radio follow-up of unidentified \fermi\ sources. The launch of \fermi\ and the extraordinary sensitivity of the Large Area Telescope \citep[LAT,][]{atwood09}, confirmed those predictions of \gray-bright MSPs with detections of eight relatively normal radio MSPs using only the first few months of \fermi\ events \citep{fermimsps}. Those MSPs were detected via the folding of \grays\ modulo the known spin and orbital ephemerides from radio timing campaigns \citep{fermiptc}. In order to best utilize radio telescope time to search either for radio counterparts to new \gray-selected pulsars or to search blindly for radio pulsations from \gray\ sources that might contain pulsars, we formed the Pulsar Search Consortium (PSC), a group of approximately 20 LAT-team members and/or pulsar experts associated with large radio telescopes around the world. This paper describes one of the PSC's first programs, which used the Green Bank Telescope (GBT) to search 25 unassociated sources from the \fermi\ LAT Bright Source List \citep{fermibsl}.

We have identified three new nearby radio MSPs as the counterparts of bright and previously unassociated {\it Fermi} LAT sources at high Galactic latitude. Our non-detection of young pulsars or MSPs in the more numerous sources searched at low Galactic latitude is likely due to our only moderate sensitivity improvements (typically 2$-$3$\times$) over the best surveys of those regions to date \citep[e.g.][]{mlc+01} due to higher sky temperatures resulting from our lower observing frequency. Additionally, the complicated and confused nature of the Galactic plane in \grays\ makes the positive identification of point sources difficult. Several of the bright sources may be blends of other sources or the result of insufficient modelling of the Galactic background. Nonetheless, deeper surveys at frequencies of 1.5$-$2\,GHz of these sources may prove more fruitful in the future. The new pulsars are very typical radio MSPs in terms of spin period, binary parameters, magnetic field strength, spin-down luminosity, and characteristic age, and their unusual brightness in \grays\ is likely due more to their proximity than to especially energetic emission processes in their magnetospheres. The very high implied \gray\ efficiency for PSR \mspa\ suggests it is likely closer, by up to a factor of 2 or more, than predicted by the NE2001 model \citep{cl02}. The line-of-sight to PSR \mspa\ is nearly tangent to the Gum Nebula where NE2001 shows an exceptionally steep DM gradient. Additionally, the pulsar's \gray\ emission is likely not isotropic, but only covers tens of percent of the sky. These large efficiencies in general, though, are consistent with the tens of percent values found by \citet{fermimsps} for radio MSPs detected in \grays\ and imply that MSPs are very efficient producers of \grays. We do not have proper motion measurements for pulsars \mspa\ or \mspb\ and so their measured spin-down rates are contaminated at some level (likely $\lesssim$10\%) by the Shklovskii effect \citep{shk70}. There is a statistically significant proper motion measurement of $\sim$100\,mas\,yr$^{-1}$ for \mspb, though, which implies a Shklovskii effect at 400\,pc larger than the measured spin-down rate for the pulsar. If the proper motion is confirmed at this level, the requirement to have the pulsar intrinsically spinning down gives an upper limit for the pulsar's distance of $\sim$240\,pc. Timing observations over the next several years will determine the proper motions and possibly the timing parallaxes for each of the pulsars. In all three cases we identified X-ray counterparts to the pulsars which substantially aided in the rapid establishment of timing solutions. The three MSPs appear to have fairly typical X-ray properties for radio MSPs \citep[e.g.][]{bgh+06} with primarily soft thermal-like spectra and X-ray luminosities in the 10$^{30-31}$\,erg\,s$^{-1}$ range, approximately 10$^{-4}$ to 10$^{-3}$ of their \gray\ luminosities. The radio flux densities of $\sim$1\,mJy near 1\,GHz are large enough to make the MSPs potentially useful for a wide variety of timing projects, such as the detection of gravitational waves via long-term pulsar timing (e.g. NANOGrav\footnote{\url{http://nanograv.org}}), yet they are small enough to explain why earlier large-area surveys for pulsars missed them \citep[e.g.][]{mld+96,lxf+05}. In addition, the fact that many of the nearby radio MSPs are being detected in \grays\ and vice-versa argues that the sizes of the radio and \gray\ beams are comparable for MSPs (likely within a factor of $\sim$2), and that deep radio and \gray\ surveys may allow us to eventually detect a large percentage of the local population of these sources. In the short-term, the fact that {\it Fermi} can point us to nearby radio MSPs is already causing a large increase in the number of known systems, with much less effort than is required by sensitive large-area radio surveys. If most radio MSPs produce \grays\ as these early results seem to indicate, MSPs may contribute to the diffuse isotropic \gray\ background \citep{fl10}.

1012.2862

1012

1012.2115_arXiv.txt

We present deep $J$-, $H$-, and $K_{s}$-band imaging data of the MOIRCS Deep Survey (MODS), which was carried out with Multi-Object Infrared Camera and Spectrograph (MOIRCS) mounted on the Subaru telescope in the GOODS-North region. The data reach 5$\sigma$ total limiting magnitudes for point sources of $J=23.9$, $H=22.8$, and $K_{s}=22.8$ (Vega magnitude) over 103 arcmin$^{2}$ (wide field). In 28 arcmin$^{2}$ of the survey area, which is ultra deep field of the MODS (deep field), the data reach the 5$\sigma$ depths of $J=24.8$, $H=23.4$, and $K_{s}=23.8$. The spatial resolutions of the combined images are FWHM $\sim 0.6$ arcsec and $\sim 0.5$ arcsec for the wide and deep fields in all bands, respectively. Combining the MODS data with the multi-wavelength public data taken with the HST, Spitzer, and other ground-based telescopes in the GOODS field, we construct a multi-wavelength photometric catalog of $K_{s}$-selected sources. Using the catalog, we present $K_{s}$-band number counts and near-infrared color distribution of the detected objects, and demonstrate some selection techniques with the NIR colors for high redshift galaxies. These data and catalog are publicly available via internet.

In recent years, several deep multi-wavelength surveys have been carried out to reveal galaxy formation and evolution in the high-redshift universe. Representative examples of such surveys are the Great Observatories Origins Deep Survey (GOODS, \cite{gia04}), the Subaru XMM-Newton Deep Survey (SXDS, K. Sekiguchi et al., in preparation), the Cosmic Evolution Survey (COSMOS, \cite{sco07}), and the All-wavelength Extended Groth strip International Survey (AEGIS, \cite{dav07}). Deep multi-wavelength observations from radio to X-ray allow us to comprehensively investigate the properties of stars, gas, dust, and AGN of high-redshift galaxies. In such multi-wavelength surveys for high-redshift galaxies, near-infrared (NIR) imaging is essential for the following reasons. First, the observed NIR luminosity of galaxies reflects their stellar mass, which is one of the most basic physical properties of galaxies, relatively well for galaxies at $z\lesssim3$. NIR data also covers the Balmer/4000\AA\ break of galaxies at $1\lesssim z \lesssim 4$, which is important to study the stellar population of these galaxies and to determine their photometric redshifts. Compared to optical light, observed NIR light is less sensitive to the effect of dust extinction, and the spatial resolution of NIR data is comparatively good even in ground-based observations. These are helpful for the identification of sources detected in other wavelengths such as X-ray, mid-IR, sub-mm and radio. The recent availability of wide-field NIR instruments with large-format detectors mounted on 4-8m telescopes has allowed us to carry out wider and deeper NIR surveys. With such new instruments, we still need a relatively long integration time in order to construct a representative sample of high-redshift galaxies. For example, a normal L$^{*}$ galaxy seen in the local universe would have $K_{\rm Vega} \sim 23 $ if placed at $z\sim3$ (e.g., \cite{fra03}). Deeper data would be desirable for the investigation of low-mass (sub-L$^{*}$) galaxies at such high redshift. The search for star-forming galaxies at $z\gtrsim7$ also requires extremely deep NIR data. Existing NIR surveys with such depth are limited to several small field surveys. They include those with the Hubble Space Telescope (HST)/NICMOS such as Hubble Deep Field North (HDF-N; \cite{dic03}; \cite{tho99}), HDF-South NICMOS field \citep{wil00}, and Hubble Ultra Deep Field (UDF, \cite{tho05}). The new HST instrument, WFC3 IR-channel is providing deeper data than those with NICMOS (e.g., \cite{win10}), but the deep NIR observations with the HST instruments are practically limited to $\lambda < 1.8 \mu$m. Although there are also ultra-deep NIR surveys with ground-based 8m class telescopes such as Subaru Deep Field \citep{mai01} and Faint Infrared Extragalactic Survey HDF-S field (FIRES; \cite{lab03}), these surveys have small fields of several arcmin$^{2}$. \begin{figure*} \begin{center} \FigureFile(140mm,140mm){f1.eps} \end{center} \vspace{-3mm} \caption{Survey field layout and exposure maps for $J$, $H$, and $K_{s}$ bands. In the top-left panel, the GOODS-N HST/ACS region and the original HDF-N region are shown in magenta and green lines. Arrows represent the direction of X- and Y-axes of the mosaiced MOIRCS images. Grayscales in the top-right, bottom-left and bottom-right panels show the integration time of the MOIRCS $J$, $H$, and $K_{s}$-band imaging, respectively. } \label{fig:fov} \end{figure*} We have carried out a ultra-deep NIR imaging survey, namely, MOIRCS Deep Survey (MODS) with Multi-Object Infrared Camera and Spectrograph (MOIRCS) mounted on the Subaru telescope in the GOODS-North region. Total integration time of $\sim$ 124 hours were spent in the $JHK_{s}$-band imaging observations with MOIRCS, and the data reach $K_{s}\sim23$ (5$\sigma$, Vega) over $\sim$ 103 arcmin$^{2}$, and $K_{s}\sim24$ over $\sim$ 28 arcmin$^{2}$. In this region, the GOODS and the Chandra Deep Field North (CDF-N) surveys provided deep multi-wavelength imaging data such as optical HST/ACS images \citep{gia04}, mid-IR images obtained with Spitzer/IRAC and MIPS (M. Dickinson et al., in preparation), and Chandra X-ray images \citep{ale03}. Extremely deep radio observations with VLA, deep sub-mm/mm surveys with SCUBA, AzTEC and MAMBO, and extensive optical spectroscopic surveys with the Keck telescopes have also been carried out (e.g., \cite{mor10}; \cite{pop06}; \cite{gre08}; \cite{per08}; \cite{bar08}; \cite{wir04}). On the other hand, NIR data which cover the GOODS-North region have reached only $m_{\rm AB} \sim $ 22--22.5 (\cite{cap04}; \cite{bun05}), although \citet{wan10} recently published relatively deep $K_{s}$-band data which reach $K_{s}\sim22.7$ (Vega). For $J$ band, there had been no wide-field data which cover most of the GOODS-N region. Therefore NIR data with a comparable depth have been desirable in this field. In this context, with the obtained ultra-deep NIR data, we have investigated the number counts of Distant Red Galaxies at $z\gtrsim2$ \citep{kaj06}, the clustering properties of stellar mass-selected sample at $1<z<4$ \citep{ich07}, the stellar mass dependence of the X-ray properties of galaxies at $2<z<4$ \citep{yam09}, the evolution of the galaxy stellar mass function at $0.5<z<3.5$ \citep{kaj09}, the correlation between the stellar mass and surface brightness for galaxies at $0.3<z<3$ \citep{ich10}, the relation between the NIR morphology and star formation activity of galaxies at $0.8<z<1.2$ \citep{kon10}, the star formation activity as a function of stellar mass at $0.5<z<3.5$ \citep{kaj10}, and the evolution of quiescent galaxies as a function of stellar mass at $0.5<z<2.5$ \citep{kaj11}. \citet{yos10} also used these NIR images with NIR spectroscopic data of star-forming BzK galaxies at $z\sim2$ to investigate the star formation activity and stellar population of these galaxies. \begin{table*} \caption{Summary of the MODS observations} \label{tab:first} \begin{center} \begin{tabular}{lccccc} \hline & center position & & exposure time & FWHM (ch1, ch2) & 5$\sigma$ limit (ch1, ch2)\footnotemark[$*$]\\ field & RA\hspace{1cm} Dec & band & (hour) & (arcsec) & (Vega mag)\\ \hline GT-1 & 12:36:24.9 +62:10:43 & $J$ & 8.0 & 0.59\hspace{5mm} 0.59 & 24.3 \hspace{5mm} 24.2 \\ & & $H$ & 2.5 & 0.58\hspace{5mm} 0.59 & 23.3 \hspace{5mm} 23.1 \\ & & $K_{s}$ & 8.3 & 0.58\hspace{5mm} 0.53 & 23.1 \hspace{5mm} 23.2 \\ \hline GT-2 & 12:36:47.8 +62:13:11 & $J$ & 28.2 & 0.48\hspace{5mm} 0.49 & 25.2 \hspace{5mm} 25.2 \\ & & $H$ & 5.7 & 0.46\hspace{5mm} 0.46 & 23.8 \hspace{5mm} 23.7 \\ & & $K_{s}$ & 28.0 & 0.45\hspace{5mm} 0.46 & 24.1 \hspace{5mm} 24.1 \\ \hline GT-3 & 12:37:09.7 +62:15:58 & $J$ & 6.3 & 0.57\hspace{5mm} 0.58 & 24.4 \hspace{5mm} 24.3 \\ & & $H$ & 3.2 & 0.55\hspace{5mm} 0.55 & 23.1 \hspace{5mm} 23.1 \\ & & $K_{s}$ & 10.7 & 0.59\hspace{5mm} 0.60 & 23.2 \hspace{5mm} 23.2 \\ \hline GT-4 & 12:37:31.8 +62:18:29 & $J$ & 9.1 & 0.58\hspace{5mm} 0.59 & 24.3 \hspace{5mm} 24.3 \\ & & $H$ & 4.3 & 0.58\hspace{5mm} 0.59 & 23.3 \hspace{5mm} 23.2 \\ & & $K_{s}$ & 9.8 & 0.59\hspace{5mm} 0.60 & 23.1 \hspace{5mm} 23.1 \\ \hline \multicolumn{6}{@{}l@{}}{\hbox to 0pt{\parbox{160mm}{\footnotesize \vspace{1mm} \par\noindent \footnotemark[$*$] the limiting magnitude is estimated from the background fluctuation measured with a aperture diameter of 2 $\times$ FWHM of the PSF. The total limiting magnitude for point sources, which is corrected for fluxes missed from the aperture, is $\sim$ 0.3 mag brighter for each field and band (see text for details). }\hss}} \label{tab:field} \end{tabular} \end{center} \end{table*} In this paper, we present the NIR imaging data and a multi-wavelength photometric catalog of $K_{s}$-selected sources. Section 2 describes the observations. We give details of the data reduction in Section 3, and the properties of the reduced data in Section 4. In Section 5, we construct the source catalog and explain the catalog entries. Using the source catalog, we present the $K_{s}$-band number counts and distribution of NIR colors in Section 6. A summary is presented in Section 7. The Vega-referred magnitude system is used throughout this paper, unless stated otherwise.

In this paper, we presented the deep $JHK_{s}$-band imaging data of the MODS obtained with MOIRCS on the Subaru telescope. The data cover an area of 103.3 arcmin$^{2}$ in the GOODS-North region and reach the 5$\sigma$ total limiting magnitudes for point sources of $J=23.9$, $H=22.8$, and $K_{s}=22.8$. In 28.2 arcmin$^{2}$ of the survey area, the data reach the 5$\sigma$ depths of $J=24.8$, $H=23.4$, $K_{s}=23.8$. The World Coordinate System of the reduced images is based on the public HST/ACS ver. 2.0 data of the GOODS. The image quality of the combined images is characterized by a PSF with FWHM of $\sim $ 0.53--0.60 arcsec for the wide field and $\sim$ 0.45--0.49 arcsec for the deep field, respectively. For the color measurements, we also provided the PSF-matched mosaic images whose FWHM of the PSFs are 0.6 arcsec for the wide field and 0.5 arcsec for the deep field. The $K_{s}$-band detection completeness for the point sources is $\sim$ 90 \% at $K_{s}\sim23$ for the wide field and at $K_{s}\sim24$ for the deep field. The spurious sources are also negligible down to these completeness limits, although the false detection rate increases rapidly below the limits. Combining the multi-wavelength public data taken with the HST, Spitzer, and other ground-based telescopes in the GOODS field with the MODS data, we constructed the multi-wavelength photometric catalogs of $K_{s}$-selected sources. Total 9875 and 3787 objects are listed in the wide and deep catalogs. The catalogs also include the Spitzer/MIPS 24 $\mu$m fluxes, Chandra X-ray fluxes, and redshifts of the $K_{s}$-selected objects. The comparisons between the spectroscopic and photometric redshifts suggest that the photometric redshift accuracy is $\delta z/(1+z_{\rm spec}) \lesssim 0.1$ with $\sim$ 3\% outliers for galaxies with spectroscopic redshifts ($\sim $ 4--12\% outliers for objects with $z_{\rm spec} > 1.5$). Using the catalogs, we examined the $K_{s}$-band number counts and NIR color distribution. The $K_{s}$-band number counts in the MODS field are consistent with those in other general fields and show the logarithmic slope of $d(\log{N})/dm \sim 0.2$ at $K\gtrsim 20$. The NIR color distribution of spectroscopically confirmed stars indicates that the zero points of the MODS data have been correctly determined. We also demonstrated some selection techniques for high-redshift galaxies with the NIR colors. The MODS data are one of the deepest NIR imaging data over a area of $\sim$ 30 to $\sim$ 100 arcmin$^{2}$ to date, especially in the $K_{s}$ band. These data sample an important part of the SEDs of galaxies at $1\lesssim z \lesssim 4$ to study their stellar population. The high image quality of the data is very useful for the deconvolution of other imaging data with poorer spatial resolution, while their depth in the NIR is essential for the identification of objects detected in the multi-wavelength data, which have been obtained by the Chandra, Spitzer, Herschel, VLA, and so on. These imaging data and $K_{s}$-selected source catalogs are publicly available on the MODS web site (http://www.astr.tohoku.ac.jp/MODS/). \bigskip We thank an anonymous referee for very helpful suggestions and comments. We also thank Kevin Bundy for kindly providing their KGOODS-N v1.0 catalog. This study is based on data collected at Subaru Telescope, which is operated by the National Astronomical Observatory of Japan. This work is based in part on observations made with the Spitzer Space Telescope, which is operated by the Jet Propulsion Laboratory, California Institute of Technology under a contract with NASA. Some of the data presented in this paper were obtained from the Multi-mission Archive at the Space Telescope Science Institute (MAST). STScI is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555. Support for MAST for non-HST data is provided by the NASA Office of Space Science via grant NAG5-7584 and by other grants and contracts. Data reduction and analysis were carried out on common use data analysis computer system at the Astronomy Data Center, ADC, of the National Astronomical Observatory of Japan. IRAF is distributed by the National Optical Astronomy Observatories, which are operated by the Association of Universities for Research in Astronomy, Inc., under cooperative agreement with the National Science Foundation. \appendix

1012.2115

1012

1012.3922_arXiv.txt

We present an overview of precise pulsar timing using data from the Large Area Telescope (LAT) on \textit{Fermi}. We describe the analysis techniques including a maximum likelihood method for determining pulse times of arrival from unbinned photon data. In addition to determining the spindown behavior of the pulsars and detecting glitches and timing noise, such timing analyses allow the precise determination of the pulsar position, thus enabling detailed multiwavelength follow up.

The \textit{Fermi} Large Area Telescope (LAT) has proven to be an exceptionally powerful instrument for the study of gamma-ray pulsars (e.g. Romani et al. in this volume). Because of the large effective area, excellent background rejection made possible by the fine point spread function (PSF), precise absolute time tagging of events, and the sky survey mode of operation, the LAT is able to do precise timing of pulsars in the gamma-ray band in a way never before possible. The continuous sky survey enables evenly sampled timing points for every pulsar in the sky over the full mission lifetime. Exploiting the LAT as a pulsar timing instrument is a requirement for the large number of pulsars that have been found in gamma-ray blind searches. Of the 26 known (Geminga plus the 25 discovered using the LAT; see Saz Parkinson, this volume), only 3 have been observed to pulse in radio wavelengths, and only one of those (PSR J2021+4127) is bright enough for routine radio timing. In addition, there are other pulsars that are more suitable for timing with the LAT than with radio observations, such as the very faint young pulsar PSR J1124$-$5916. Lastly, some very bright gamma-ray pulsars, such as the Crab and Vela, are good targets for LAT timing because they can be timed precisely in the gamma-ray band without any concern from time-variable dispersion measure or other propagation effects that can afflict the radio observations. Pulsar timing allows us to characterize the rotational parameters of the pulsar and study the effects of timing noise and glitches, but for the newly-discovered gamma-ray selected pulsars perhaps the most important measurement enabled by LAT timing is precise position determination. These timing positions achieve arcsecond accuracy, much better than the several arcminute accuracies that come from LAT photon direction measurements, and these enable deep searches for multiwavelength counterparts.

As we have shown, the \textit{Fermi} LAT is a powerful instrument for gamma-ray pulsar timing. LAT timing is an important tool for studying a range of pulsars from the very brightest to those that have yet to be detected in any other wave band. LAT TOAs achieve sub-millisecond accuracy on many pulsars and tens of microsecond accuracy on bright millisecond pulsars and the very bright Crab and Vela pulsars. LAT timing models have been determined for all of the pulsars discovered in blind searches of the LAT data as well as a number of pulsars that are better timed with the LAT. An important result of the timing analysis is the determination of precise (arcsecond) positions for these pulsars that enable multiwavelength follow up in the radio, optical, and X-ray bands. Timing parameters for LAT pulsars are made available via the \textit{Fermi} Science Support Center\footnote{\url{http://fermi.gsfc.nasa.gov/ssc/data/access/lat/ephems/}}. \begin{theacknowledgments} The \textit{Fermi} LAT Collaboration acknowledges support from a number of agencies and institutes for both development and the operation of the LAT as well as scientific data analysis. These include NASA and DOE in the United States, CEA/Irfu and IN2P3/CNRS in France, ASI and INFN in Italy, MEXT, KEK, and JAXA in Japan, and the K.~A.~Wallenberg Foundation, the Swedish Research Council and the National Space Board in Sweden. Additional support from INAF in Italy and CNES in France for science analysis during the operations phase is also gratefully acknowledged. \end{theacknowledgments}

1012.3922

1012

1012.3749_arXiv.txt

High resolution, multi-wavelength maps of a sizeable set of nearby galaxies have made it possible to study how the surface densities of \hi, \htwo\ and star formation rate ($\Sigma_{\rm HI}$,$\Sigma_{\rm H2}$,$\Sigma_{\rm SFR}$) relate on scales of a few hundred parsecs. At these scales, individual galaxy disks are comfortably resolved, making it possible to assess gas-SFR relations with respect to environment within galaxies. $\Sigma_{\rm H2}$, traced by CO intensity, shows a strong correlation with $\Sigma_{\rm SFR}$ and the ratio between these two quantities, the molecular gas depletion time, appears to be constant at about 2\,Gyr in large spiral galaxies. Within the star-forming disks of galaxies, $\Sigma_{\rm SFR}$ shows almost no correlation with $\Sigma_{\rm HI}$. In the outer parts of galaxies, however, $\Sigma_{\rm SFR}$ does scale with $\Sigma_{\rm HI}$, though with large scatter. Combining data from these different environments yields a distribution with multiple regimes in $\Sigma_{\rm gas}-\Sigma_{\rm SFR}$ space. If the underlying assumptions to convert observables to physical quantities are matched, even combined datasets based on different SFR tracers, methodologies and spatial scales occupy a well define locus in $\Sigma_{\rm gas}-\Sigma_{\rm SFR}$ space.

Great progress has been made towards an understanding of star formation (SF) on small scales in the Milky Way, but many open questions remain about its connection to large scale processes: what sets where SF occurs in galaxies and how efficiently gas is converted into stars? How important are global, galaxy-scale environmental parameters as opposed to small-scale properties of the interstellar medium (ISM)? What is the role of feedback in regulating SF? To address such questions, theoretical modeling and simulations need to be constrained by comprehensive observations. Both observations and theory have focused on the relationship between the star formation rate (SFR) and the gas density, for which a tight power-law relationship was observed in a large number of galaxies by \cite{kennicutt98}. Such a relationship was first suggested many decades ago by \cite{schmidt59}, who studied the distributions of atomic gas and stars in the Galaxy. Over the following decades, similar studies targeted individual Local Group galaxies, e.g., M33 (\cite{madore74,newton80}), the Large Magellanic Cloud (\cite{tosa75}), and the Small Magellanic Cloud (\cite{sanduleak69}). \cite{kennicutt89} carried out the first comprehensive extragalactic study targeting a large sample of nearby galaxies and \cite{kennicutt98} followed up this work, focusing on measurements averaged across galaxy disks. In a sample of 97 nearby normal and starburst galaxies, he found a close correlation between the galaxy-average total gas surface density ($\Sigma_{\rm gas} = \Sigma_{\rm HI} + \Sigma_{\rm H2}$) and the galaxy-average SFR surface density ($\Sigma_{\rm SFR}$). Following this work, it has become standard to study the relationship between gas and star formation via surface densities, which are observationally more easily accessible than volume densities. \cite{kennicutt98} found $\Sigma_{\rm SFR}=A\times\Sigma_{\rm gas}^{N}$, with intercept $A$ and power law index $N$ --- a relationship that is variously referred to as the ``star formation law,'' ``Schmidt-Kennicutt law,'' or ``Schmidt Law.'' \cite{kennicutt98} derived $N\approx1.40$. Because the ratio $\Sigma_{\rm SFR}/\Sigma_{\rm gas}$ describes how efficiently gas is converted into stars (and is thus often referred to as the star formation efficiency, SFE), this super-linear power law index implies that systems with higher average gas surface densities more efficiently convert gas into stars (left panel, Figure \ref{fig1}). This measured value is close to $N=1.5$, which is expected if the free-fall time in a fixed scale height gas disk is the governing timescale for SF on large scales. Other studies working with disk-averaged, global measurements found $N$ to be in the range of $\sim0.9-1.7$ (e.g., \cite{buat89,buat92,deharveng94}). \begin{figure}[t] \begin{center} \includegraphics[width=1.5in]{bigiel-fig1-1.eps} \includegraphics[width=1.5in]{bigiel-fig1-2.eps} \includegraphics[width=1.5in]{bigiel-fig1-3.eps} \caption{{\em Left:} \cite{kennicutt98} found a strong correlation over many orders-of-magnitude between global averages of $\Sigma_{\rm SFR}$ and $\Sigma_{\rm gas} = \Sigma_{\rm HI}+\Sigma_{\rm H2}$ for a large sample of nearby normal and starburst galaxies. He derived a power law index $N\approx1.40$, implying more efficient SF for galaxies with higher average gas columns. {\em Middle and Right:} Surface densities of gas and SFR measured in radial profiles by \cite{wong02} for two exemplary nearby spirals. The middle panel shows $\Sigma_{\rm SFR}$ versus $\Sigma_{\rm HI}$, the right panel $\Sigma_{\rm SFR}$ versus $\Sigma_{\rm H2}$. \cite{wong02} found no correlation between atomic gas and SFR, whereas molecular gas and SFR scale with one another.} \label{fig1} \end{center} \end{figure} The availability of high-resolution maps of CO emission, the standard tracer of molecular gas, made it possible to follow up the work of \cite{kennicutt98} with studies focusing on azimuthally-averaged radial profiles of gas and SF. Resolving galaxies in this way makes it possible to look at how gas and SF relate within individual galaxy disks, opening up a wide range of environmental factors to explore. \cite{wong02} used BIMA SONG data (\cite{helfer03}) to study 6 nearby spirals, \cite{boissier03} explored a larger sample of nearby spirals, \cite{heyer04} studied the Local Group galaxy M33, and \cite{schuster07} explored the gas-SF relation in M51. These studies derived power law indices in the range $N\approx1-3$, leaving it unclear whether a single relation relates gas and SF when galaxy disks are spatially resolved. Further disagreement centered on the relationship of SF to different types of gas --- \hi , H$_2$, and total gas. Intuitively, one might expect a stronger correlation between SF and the cold, molecular phase, rather than the atomic phase. \cite{wong02} indeed found a much stronger correlation of $\Sigma_{\rm SFR}$ with the molecular gas, $\Sigma_{\rm H2}$ (compare Figure \ref{fig1}). However, \cite{kennicutt98} and \cite{schuster07} both found a better correlation of $\Sigma_{\rm SFR}$ with the total gas, $\Sigma_{\rm gas}$, than with $\Sigma_{\rm H2}$.

With vast improvements in the data available for nearby galaxies some consensus is beginning to emerge on how different parts of galaxies populate the $\Sigma_{\rm SFR}$-$\Sigma_{\rm gas}$ parameter space. The role of environmental quantities other than gas surface density alone are beginning to become clear and different relations are emerging for different types and parts of galaxies. When viewed in detail the composite relation may not be a simple power law, but it contains key information to constrain theories and to benchmark simulations. Challenges remain, too. The determination of star formation rates at low surface brightness is still difficult. The use of CO to trace \htwo\ underpins almost all of this work but the CO-to-\htwo\ conversion factor remains imprecisely calibrated as a function of environment. Finally, the fundamental units of star-formation, individual molecular clouds, remain largely observationally inaccessible beyond the Local Group --- a situation that will not change until ALMA begins its full operations.

1012.3749

1012

1012.1875_arXiv.txt

Very luminous stars, with bolometric luminosities exceeding $10^6 L_\odot$, appear spectroscopically mainly as Wolf-Rayet (WR) types. However, hot stars emit most of their radiation in the extreme ultraviolet which is not accessible to observation. Hence the determination of their luminosity $L$ must rely on adequate model atmospheres. Our ``Potsdam Wolf-Rayet'' code (PoWR -- see Hamann \& Gr\"afener 2003 and references therein) solves the non-LTE radiative transfer in a spherically symmetric expanding atmosphere. Detailed and complex model atoms are taken into account especially for H, He, and the CNO elements, while the iron-group elements are treated in the superlevel approximation. Wind inhomogeneities are accounted for in a first-order approximation (``microclumping''). The code has been applied mainly for the wind-dominated emission-line spectra of WR stars (see {\tt http://www.astro.physik.uni-potsdam.de/PoWR.html} for grids of models), but can also be used for fitting photospheric absorption spectra. In the standard version of the PoWR code, mass-loss rate and velocity field are free parameters of the model, while they are determined consistently with the radiation pressure only in the hydrodynamical version (Gr\"afener \& Hamann 2005, 2008). Another PoWR code option not used here is ``macroclumping'' (Oskinova et al.\ 2007).

1012.1875

1012

1012.4457_arXiv.txt

We present a comprehensive survey of boron abundances in diffuse interstellar clouds from observations made with the Space Telescope Imaging Spectrograph (STIS) of the \emph{Hubble Space Telescope}. Our sample of 56 Galactic sight lines is the result of a complete search of archival STIS data for the B~{\small II} $\lambda$1362 resonance line, with each detection confirmed by the presence of absorption from O~{\small I}~$\lambda$1355, Cu~{\small II}~$\lambda$1358, and Ga~{\small II}~$\lambda$1414 (when available) at the same velocity. Five previous measurements of interstellar B~{\small II} from Goddard High Resolution Spectrograph observations are incorporated in our analysis, yielding a combined sample that more than quadruples the number of sight lines with significant boron detections. Our survey also constitutes the first extensive analysis of interstellar gallium from STIS spectra and expands on previously published results for oxygen and copper. The observations probe both high and low-density diffuse environments, allowing the density-dependent effects of interstellar depletion to be clearly identified in the gas-phase abundance data for each element. In the case of boron, the increase in relative depletion with line-of-sight density amounts to an abundance difference of 0.8~dex between the warm and cold phases of the diffuse interstellar medium. The abundance of boron in warm, low-density gas is found to be B/H~=~$(2.4\pm0.6)\times10^{-10}$, which represents a depletion of 60\% relative to the meteoritic boron abundance. Beyond the effects of depletion, our survey reveals sight lines with enhanced boron abundances that potentially trace the recent production of $^{11}$B, resulting from spallation reactions involving either cosmic rays or neutrinos. Future observations will help to disentangle the relative contributions from the two spallation channels for $^{11}$B synthesis.

The origins of the two stable isotopes of boron, $^{10}$B and $^{11}$B, remain uncertain despite numerous theoretical and observational advances over the past four decades (see the reviews by Reeves 1994; Vangioni-Flam et al. 2000; Prantzos 2007). As one of the rare light elements (a group that also includes lithium and beryllium), boron cannot be synthesized through quiescent burning in stellar interiors, where it is destroyed at temperatures that exceed $5\times10^6$ K, nor is it produced in significant quantities by the standard model of the Big Bang. The spallation of interstellar nuclei by relativistic Galactic cosmic rays (GCR), as originally proposed by Reeves et al. (1970), is known to be an effective means of producing the light elements. However, detailed models of GCR spallation (e.g., Meneguzzi et al. 1971, hereafter MAR; Ramaty et al. 1997), while successful at reproducing the solar system abundance of $^{10}$B (as well as $^9$Be and, by including the $\alpha$-$\alpha$ fusion reactions, $^6$Li), fail to adequately account for $^{11}$B (and $^7$Li). Standard GCR nucleosynthesis predicts a boron isotopic ratio of $^{11}$B/$^{10}$B~=~2.4 (MAR), in conflict with the value measured for solar system material (4.0; Lodders 2003). Still, since the threshold energies for spallation reactions leading to $^{11}$B are generally lower than those for reactions that produce $^{10}$B, it is possible to account for the discrepancy by adopting a cosmic-ray spectrum enhanced at low energies (particularly in the range 5$-$40~MeV nucleon$^{-1}$; see MAR; Meneguzzi \& Reeves 1975). The GCR energy spectrum below 1 GeV nucleon$^{-1}$ is highly uncertain due to the effects of solar modulation, meaning that there are no direct observational constraints against a large flux of low-energy cosmic rays. In the present-day interstellar medium (ISM), the dominant contribution to cosmic-ray nucleosynthesis comes from forward spallation reactions (i.e., energetic protons and $\alpha$-particles impinging on interstellar CNO nuclei). Forward spallation is a secondary production mechanism because it depends on the metallicity of the ISM (i.e., the CNO abundances) and on the rate of supernovae occurring in the Galaxy (supernovae being the presumed sources of GCR acceleration). Over Galactic evolutionary timescales, the abundances of light elements synthesized via forward spallation should scale quadratically with the abundances of the metals that serve as interstellar targets. However, this fact is contrary to the well-known primary behavior exhibited by beryllium and boron in the halo. Studies of metal-poor halo stars (e.g., Gilmore et al. 1992; Duncan et al. 1992, 1997; Garc\'ia L\'opez et al. 1998; Boesgaard et al. 1999) have consistently shown that Be and B abundances increase approximately linearly with metallicity. Thus, many investigations (e.g., Cass\'e et al. 1995; Ramaty et al. 1996; Vangioni-Flam et al. 1996) have focused on scenarios involving reverse spallation reactions (i.e., accelerated CNO nuclei being spalled from ambient interstellar H and He). In these models, low-energy, metal-enriched cosmic rays are accelerated in superbubbles (Parizot \& Drury 1999, 2000; Parizot 2000) by shocks associated with Type II supernovae (SNe II) or by the winds of massive Wolf-Rayet stars, and then interact with nearby interstellar material. Since light elements produced in superbubbles result from the breakup of freshly synthesized CNO nuclei, the superbubble model represents a primary mechanism that can operate throughout the lifetime of the Galaxy and may be particularly important at early times. Neutrino-induced spallation in SNe II (the $\nu$-process; Woosley et al. 1990) offers an alternative explanation for primary boron production and could ameliorate the problem in standard GCR nucleosynthesis of a low predicted $^{11}$B/$^{10}$B ratio. A significant amount of $^{11}$B (though virtually no $^{10}$B) is expected to be synthesized during the collapse of a massive star's core as the immense flux of neutrinos interacts with $^{12}$C in the carbon-rich shell. Models of core-collapse supernovae that incorporate the $\nu$-process (see Woosley \& Weaver 1995, hereafter WW95) also predict substantial yields for $^7$Li (produced in the helium shell) and $^{19}$F (in the neon shell). Timmes et al. (1995), adopting the WW95 yields for their model of Galactic chemical evolution (GCE), conclude that a major portion of the cosmic abundances of $^{11}$B and $^{19}$F, and about half of the $^7$Li abundance, can be attributed to neutrino nucleosynthesis. However, these prescriptions may be difficult to accommodate if other sources [e.g., cosmic-ray spallation for $^{11}$B and $^7$Li, Big Bang nucleosynthesis for $^7$Li, stellar processing in asymptotic giant branch (AGB) stars for $^7$Li and $^{19}$F] must also contribute. In the spallation models of Fields et al. (2000), which include a GCR, superbubble, and neutrino component, the $\nu$-process yields of WW95 had to be reduced by 60\% in order to avoid the overproduction of $^{11}$B. The $\nu$-process interactions chiefly involve the $\mu$- and $\tau$-neutrinos because these have higher temperatures, yet the temperatures are uncertain and the light element yields are strongly dependent on these values. Yoshida et al. (2005, 2008) constrained the $\mu$- and $\tau$-neutrino temperatures by requiring that $\nu$-process synthesis not result in $^{11}$B overproduction, in accordance with GCE models that include cosmic-ray spallation (e.g., Ramaty et al. 2000; Fields et al. 2000). The constraints place the temperatures between 60\% and 80\% of the value adopted in WW95, though even lower temperatures are inferred if neutrino oscillation effects are taken into account (see Yoshida et al. 2006, 2008). The extent to which the various nucleosynthetic processes contribute to the enhancement in the solar system abundance of $^{11}$B over that predicted by standard GCR spallation is still unclear. However, since all of the boron production mechanisms occur in, or are closely associated with, the ISM, additional clues could potentially be gleaned from a careful study of interstellar boron abundances and isotopic ratios. Boron was first detected in the ISM by Meneguzzi \& York (1980), who used the \emph{Copernicus} satellite to measure absorption from the B~{\small II}~$\lambda$1362 resonance line toward $\kappa$~Ori. They obtained an interstellar abundance (B/H = $1.5\times10^{-10}$) in good agreement with the then-current stellar value ($2\times10^{-10}$; Boesgaard \& Heacox 1978), assumed to be the Galactic value. Later investigations with the Goddard High Resolution Spectrograph (GHRS) on board the \emph{Hubble Space Telescope} (\emph{HST}; e.g., Federman et al. 1993; Jura et al. 1996) added new sight lines to the list of detections (all toward nearby stars) and found that the interstellar abundances were substantially lower than both the solar and stellar values. This led Jura et al. (1996) to propose a scenario involving either the recent infall of metal-poor material in the vicinity of the Sun or the depletion of boron onto interstellar grains. Federman et al. (1996a), examining GHRS data for the line of sight to $\delta$~Sco, provided the first measurement of the $^{11}$B/$^{10}$B ratio outside the solar system. Their result ($^{11}$B/$^{10}$B = 3.4$^{+1.3}_{-0.6}$), along with the subsequent work of Lambert et al. (1998), showed that the solar system ratio is not anomalous but probably representative of the local Galactic neighborhood. The B~{\small II} survey by Howk et al. (2000) expanded the sample of interstellar boron abundances to include some of the more extended sight lines accessible to the Space Telescope Imaging Spectrograph (STIS) of \emph{HST}. These authors found that the gas-phase B/O ratio decreases with increasing average hydrogen density and interpreted the trend as an indication of interstellar depletion. From the least depleted sight line, they derived a lower limit to the present-day total interstellar boron abundance of B/H $\gtrsim$ $(2.5\pm0.9)\times10^{-10}$. The discovery by Knauth et al. (2000) of newly synthesized lithium toward $o$ Per, a member of the Per OB2 association, was an important new development in the study of light element nucleosynthesis. At least one of the clouds in the $o$ Per direction has a low $^7$Li/$^6$Li ratio ($\sim2$; Knauth et al. 2000, 2003b), consistent with the ratio predicted by standard GCR spallation (1.5; MAR) and much lower than the solar system value (12.2; Lodders 2003). The line of sight to $o$ Per passes very near the massive star-forming region IC~348 and measurements of interstellar OH suggest that the sight line possesses an order-of-magnitude higher cosmic-ray flux compared to other sight lines in Per OB2 (Federman et al. 1996b). The implication is that accelerated particles supplied by the star-forming region are interacting with ambient interstellar material in the direction of $o$ Per, leading to an enhancement in $^6$Li relative to $^7$Li. This unexpected result prompted our STIS program (GO 8622) aimed at measuring $^{11}$B/$^{10}$B ratios in Per OB2 from observations of the B~{\small II} line toward four stars: 40 Per, $o$ Per, $\zeta$ Per, and X Per. Ultimately, the acquired STIS spectra lacked the signal-to-noise (S/N) ratio necessary to yield meaningful results on $^{11}$B/$^{10}$B, though the data did provide accurate B~{\small II} column densities. Therefore, we redirected our efforts to obtaining elemental boron abundances for a larger, more statistically significant sample of Galactic sight lines. Such a sample is needed to determine conclusively the level of boron depletion in interstellar gas. This information would then enable a more robust interpretation of the observed abundances and possibly allow the detection of intrinsic abundance variations, which, if discovered, could offer vital clues to the nucleosynthetic origin of boron and other light elements. In this investigation, we more than quadruple the number of interstellar sight lines with significant detections of the B~{\small II} line by taking advantage of the wealth of UV data provided by the \emph{HST}/STIS archive. The remainder of this paper is organized as follows. We describe the observations and our reduction of the data in \S{}~2. The methods used to obtain column densities (i.e., the integration of apparent column density profiles and the method of profile synthesis) are detailed in \S{}~3, where we also provide a comparison with previous studies. The profile synthesis results are utilized in \S{}~4 to determine elemental abundances (\S{}~4.1) and depletions (\S{}~4.2) and trends are sought with various measures of gas density. In \S{}~4.3, the boron abundances are examined in detail to search for intrinsic variations superimposed on the general trend due to depletion. The results of our analyses are discussed in \S{}~5 in the context of light element nucleosynthesis and the summary and conclusions, along with suggestions for future studies, are given in \S{}~6.

Archival STIS spectra sampling diffuse interstellar gas along over 50 Galactic sight lines were analyzed to obtain column densities of B~{\small II}, and, in the process, O~{\small I}, Cu~{\small II}, and Ga~{\small II}. Before synthesizing the B~{\small II} absorption profile along a given line of sight, a consistent component decomposition was derived in the stronger UV lines. The component structure for many of these sight lines is confirmed through the analysis of high resolution Ca~{\small II} and K~{\small I} profiles. Both UV and visible data were used to construct profile templates that were fit to the B~{\small II} line, yielding the total column density along the line of sight. For sight lines with multiple complexes of absorption components, separate templates were fit to each complex, independently. This procedure allowed the detection of a slight increase in the boron-to-oxygen ratio in the inner Sagittarius-Carina spiral arm, a result which may indicate that boron production in the current epoch is dominated by a secondary process. Elemental abundances in the gas phase were determined for each of the UV species and were compared with line-of-sight measures of gas density to quantify the effects of depletion onto interstellar grains. The depletion in each element was found to increase with the average density of hydrogen, which characterizes the overall concentration of cold clouds relative to warm gas along the line of sight. Mean abundances were determined for the warm and cold phases of the diffuse ISM and the level of depletion in warm gas was shown to increase with the condensation temperature of the element. The gas-phase abundance of boron in the warm diffuse ISM was found to be B/H = $(2.4\pm0.6)\times10^{-10}$, which translates into a depletion of 60\% relative to the meteoritic boron abundance. Knowledge of the trend of decreasing elemental abundance with increasing gas density allowed the identification of sight lines showing enhanced boron abundances. Many of these sight lines (e.g., HD~93222, HD~114886, HD~208947, HD~148937, HD~43818, and HD~203374) are near regions of massive star formation, and therefore (at least some of them) may be probing the recent production of $^{11}$B, resulting from spallation reactions induced by either cosmic rays or neutrinos. Further detailed analysis is needed to determine conclusively the degree of ionization and depletion in the gas in these directions. Observations of other light elements will also be required to disentagle the relative contributions from the two spallation channels, if recent nucleosynthesis can be confirmed. As with the majority of sight lines in our survey, the complicated component structure in most of the directions showing boron enhancements will not permit any meaningful determinations of the $^{11}$B/$^{10}$B ratio. Additional measurements on interstellar fluorine may help to clarify the situation if abundances can be obtained for a larger number of sight lines from the boron sample. The elemental F/B ratio is expected to be an important tracer of neutrino spallation, a process for which direct observational evidence is still lacking. From a limited sample, we find no indication that the $\nu$-process has operated in the Cepheus bubble (as probed by the line of sight to HD~209339), though this region is known to have been shaped by core-collapse supernovae in the recent past. Future searches for observational signatures of light element production in the Galaxy should focus on specific regions where active nucleosynthesis is likely to be occurring. Among the more promising candidates are the interstellar clouds in the vicinity of the interacting supernova remnant IC~443 and those near the massive star-forming region IC~348. Observations of Li~{\small I} along lines of sight through the supernova remnant are currently being analyzed and similar data exist on Li~{\small I} toward members of IC 348. Complementary observations of B~{\small II} toward stars in both of these regions should now be acquired so that a more complete picture of light element synthesis can be obtained. Such observations have recently become feasible with the installation of the Cosmic Origins Spectrograph (COS) on \emph{HST}. In conjunction with future studies in the Galactic ISM, it will be important to expand the investigation of interstellar boron to the Large and Small Magellanic Clouds. Interstellar lithium has very recently been detected in the Small Magellanic Cloud (Howk 2010) and boron could likely be discovered there with COS or even STIS, now that it has been refurbished. The Magellanic Clouds present a unique opportunity to study light element nucleosynthesis in metal-poor environments with regions of active star formation, providing suitable analogs of our own Galaxy at earlier times. A similarly important discovery will be the eventual detection of beryllium in interstellar space, as the sole stable isotope of beryllium ($^9$Be) can be produced only through cosmic-ray spallation. While existing searches for Be~{\small II} have mainly focused on bright interstellar targets like $\zeta$~Per, $\zeta$~Oph, and $\delta$~Sco (see H\'ebrard et al. 1997), beryllium will be more heavily depleted in these directions if it follows a trend similar to that found for boron. Absorption from Be~{\small II} should be sought in more diffuse directions where less depletion is expected, though even in favorable conditions the equivalent widths will likely be less than $\sim0.1$ m\AA. Still, with sufficiently long exposure times, a detection should be possible considering the capabilities of modern large aperture telescopes. The near future promises considerable advancements in the field of cosmochemistry. For instance, actual measurements of the flux of low-energy cosmic rays, needed for accurate predictions of light element synthesis via cosmic-ray spallation, are on the horizon. The \emph{Voyager 1} and \emph{2} spacecraft, which recently crossed the solar wind termination shock, may be able to measure the flux of cosmic rays at energies below 1 GeV nucleon$^{-1}$ when they reach the heliopause in the coming decades. These measurements could help answer exciting new questions about whether anomalous cosmic rays accelerated in astrospheres throughout the Galaxy contribute significantly to the GCR flux at low energies, an outcome with direct implications for light element nucleosynthesis. The comprehensive survey of interstellar boron presented here, when combined with future studies of light element abundances and isotopic ratios, an improved treatment of interstellar depletion, and a direct determination of the flux of low-energy cosmic rays, will ultimately lead to a deeper understanding of the mechanisms responsible for the production of the light elements.

1012.4457

1012

1012.3331_arXiv.txt

We report on the discovery of ETHOS~1 (PN G068.1$+$11.0), the first spectroscopically confirmed planetary nebula (PN) from a survey of the SuperCOSMOS Science Archive for high-latitude PNe. ETHOS~1 stands out as one of the few PNe to have both polar outflows (jets) travelling at $120\pm10$ km/s and a close binary central star. The lightcurve observed with the Mercator telescope reveals an orbital period of 0.535 days and an extremely large amplitude (0.816 mag) due to irradiation of the companion by a very hot pre-white dwarf. ETHOS~1 further strengthens the long suspected link between binary central stars of planetary nebulae (CSPN) and jets. INT IDS and VLT FORS spectroscopy of the CSPN reveals weak N III, C III and C IV emission lines seen in other close binary CSPN and suggests many CSPN with these weak emission lines are misclassified close binaries. We present VLT FORS imaging and Manchester Echelle Spectrometer long slit observations from which a kinematic model of the nebula is built. An unusual combination of bipolar outflows and a spherical nebula conspire to produce an $X$-shaped appearance. The kinematic age of the jets ($1750\pm250$ yrs/kpc) are found to be older than the inner nebula ($900\pm100$ yrs/kpc) consistent with previous studies of similar PNe. Emission line ratios of the jets are found to be consistent with reverse-shock models for fast low-ionisation emitting regions (FLIERS) in PNe. Further large-scale surveys for close binary CSPN will be required to securely establish whether FLIERS are launched by close binaries.

In the protracted debate on the shaping mechanisms of PNe (Balick \& Frank 2002) appropriate solutions are sought to explain both the dominant nebula morphology (e.g. spherical, elliptical, bipolar; Balick 1987) and the accompanying collimated outflows or `jets'. Single star evolution readily explains spherical PNe and potentially elliptical PNe if there is interaction with the interstellar medium (ISM; Villaver et al. 2003; Wareing et al. 2007). Highly axisymmetric or bipolar nebulae (Corradi \& Schwarz 1995) have been modelled with some success including the generalised interacting stellar wind model (GISW; Kwok, Purton \& Fitzgerald 1978; Kwok 2000) and a combination of stellar rotation and magnetic fields (e.g. Garc\'ia-Segura 1997). There are however fundamental limitations in both of these models. The GISW model depends upon an \emph{assumed} density contrast to produce bipolar PNe (e.g. a dusty torus) and the latter cannot operate without an additional supply of angular momentum (e.g. a binary companion; Soker 2006; Nordhaus, Blackman \& Frank 2007). While there are multiple advantages to a binary explanation over existing theories, we are only beginning to understand the prevalence and impact of binarity in PNe (De Marco 2009; Miszalski et al. 2009a, 2009b, 2010). Decisive observational evidence is required to advance the shaping debate which is imbalanced towards theoretical models. PNe with dusty disks seem to support the GISW model though there appears to be a strong dependence on a binary companion for their formation (Chesneau 2010). Similarly, magnetic fields are difficult to observe directly in PNe (Jordan et al. 2005; Sabin et al. 2007), yet they appear to be more influential as part of a binary-driven dynamos (Nordhaus \& Blackman 2006; Nordhaus et al. 2007). Evidence for close binary central stars of PNe (CSPN) that passed through a common-envelope (CE) phase is now firmly established and gaining considerable momentum with at least 40 now known (Miszalski et al. 2009a, 2010; De Marco et al. 2008). These make up at least $17\pm5$\% of all CSPN (Miszalski et al. 2009a; Bond 2000). With a larger sample of 30 post-CE nebulae Miszalski et al. (2009b) found at least 30\%, perhaps 60--70\%, of post-CE binaries had bipolar nebulae suggesting the CE phase preferentially shapes bipolar nebulae. This is a substantial improvement over the study of Bond \& Livio (1990) whose sample of around a dozen close binaries showed no clear morphological trends. High resolution kinematics of nebulae around post-CE CSPN have shown a tendency for matched orbital and nebulae inclinations that strengthen the connection between the binary and its nebula (e.g. Mitchell et al. 1997; Jones et al. 2010). Further progress in this aspect will require a statistically significant sample of post-CE nebulae to be identified and studied in detail to reveal trends bestowed upon the nebula during the CE phase (Miszalski et al. 2010). Miszalski et al. (2009b) made the first steps in this direction by finding a propensity for low-ionisation structures (Gon{\c c}alves et al. 2001) amongst post-CE nebulae either in the orbital plane as a ring (e.g. Sab 41, Miszalski et al. 2009b; the Necklace, Corradi et al. 2010) and in the polar direction as jets (e.g. A~63, Mitchell et al. 2007). Jets occur in a wide variety of astrophysical objects including PNe in which they are the least understood (Livio 1997). The connection between jets and binarity has long been suspected (Soker \& Livio 1994) but never proven given the paucity of known PNe with jets \emph{and} binary CSPN. Jets in PNe typically have low outflow velocities ($\sim$100 km/s), but higher velocities of 200--300 km/s are not uncommon, e.g. M 1-16 (Schwarz 1992; Corradi \& Schwarz 1993), NGC 6337 (Corradi et al. 2000) and NGC 2392 (Gieseking, Becker \& Solf 1985), and can even reach 630 km/s in the case of MyCn 18 (O'Connor et al. 2000). There is no strong evidence for highly collimated outflows (e.g. Herbig-Haro-like outflows) in genuine PNe and Frew \& Parker (2010) suggest the few objects with the most collimated outflows are instead related to B[e] stars. The most striking example is Hen 2-90 (Sahai \& Nyman 2000) which was later reclassified as a B[e] star (Kraus et al. 2005; Frew \& Parker 2010). Instead, jets in PNe manifest as point-symmetric `corkscrew'-like outflows (L\'opez et al. 1993; V\'azquez et al. 2008) or as pairs of opposing knots often called fast low-ionisation emitting regions (FLIERS; see e.g. Balick et al. 1987, 1993; Dopita 1997) with notable examples being Hb 4 (L\'opez et al. 1997; Harrington \& Borkowski 2000; Miszalski et al. 2009b), NGC 6337 (Corradi et al. 2000; Garc\'ia-D\'iaz et al. 2009), A~63 (Mitchell et al. 2007), IC 4634 (Guerrero et al. 2008), IC 4673 (Kovacevic et al. 2010), Sab 41 (Miszalski et al. 2009b) and the Necklace (Corradi et al. 2010). Further examples are listed by Gon{\c c}alves et al. (2001) and Harrington \& Borkowski (2000). Point-symmetric jets have received the most attention in the literature with models consisting of an episodic precessing jet powered by an accretion disk in a wide binary with $P_\mathrm{orb}$$\sim$a few years (e.g. Raga, Cant\'o \& Biro 1993; Cliffe et al. 1995; Haro-Corzo et al. 2009; Raga et al. 2009). Precession helps smear out or widen the jet and may also explain the often `bent' position of jets with respect to the presumed major axis (e.g. Hb 4 and Sab 41). Some jets appear in two pairs (e.g. NGC 6337) and these may have arose from multiple ejections (e.g. Nordhaus \& Blackman 2006). A similar scenario involves a \emph{short-lived} `wobbling' accretion disk (Livio \& Pringle 1996, 1997; Icke 2003) and models with magnetic fields have also successfully been applied to the problem (Garc\'ia-Segura 1997). Dynamical lifetimes of jets suggest they are ejected well before the main nebula (Mitchell et al. 2007; Corradi et al. 2010), i.e. before the CE is ejected if a close binary is present, and this is consistent with the short-lived nature of the modelled accretion disks. The most constructive way to establish the link between jets and binaries is to look for binary CSPN in a large sample of PNe with jets. This is one of the major aims behind our ongoing program to look for close binary CSPN via the photometric monitoring technique (Miszalski et al. 2010). We are also targeting bipolar PNe and through non-detections we will also be able to discern whether larger orbital periods produce point-symmetric jets rather than FLIERS. The program has been very successful so far with at least five close binary CSPN found using the 1.2-m Flemish Mercator telescope. These include the Necklace and its Supernova 1987A-like ring (Corradi et al. 2010), the Type Ia supernova candidate double-degenerate nucleus of Hen 2-428 (Santander-Garc\'ia et al. 2010), and the subject of this paper, ETHOS~1. This paper is structured as follows. Sec. \ref{sec:discovery} introduces the new survey responsible for the discovery of ETHOS~1 and the spectroscopic confirmation of its PN status. Spectroscopic and photometric evidence for binarity are presented in Sec. \ref{sec:binary}. The nebula is examined in Sec. \ref{sec:neb} with narrow-band images, high-resolution spectroscopy. Sec. \ref{sec:diag} examines diagnostic emission line ratios of the jets and we conclude in Sec. \ref{sec:conclusion}.

\label{sec:conclusion} We have introduced ETHOS~1 (PN~G068.1$+$11.1), the first spectroscopically confirmed PN discovered from the Extremely Turquoise Halo Object Survey (ETHOS, Miszalski et al. in prep). An irradiated close binary central star was discovered with an orbital period of 0.535 days and an amplitude of 0.816 mag. The extreme amplitude is second only to Sab 41 (Miszalski et al. 2009a) and is consistent with the presence of a very hot central star that produces the highly ionised nebula ($T_e=17700$ K). The Necklace (Corradi et al. 2010) and K 1-2 (Exter et al. 2003) also share similarly large amplitudes and high-ionisation nebulae, although the absence of low-ionisation filaments in ETHOS~1 may suggest a slightly different evolutionary history. VLT FORS spectroscopy of the CSPN confirms the presence of a close binary with a large velocity separation between primary and secondary components and the nebula. The presence of N III, C III and C IV emission lines continues the trend seen in other irradiated close binary CSPN. These weak emission lines are typical of many CSPN classified as \emph{wels} in the literature and we expect many of these will turn out to be close binaries. Further observations are required to constrain the orbital inclination, masses and radii of the binary CSPN. A spectacular pair of jets travelling at $120\pm10$ km/s accompanies the inner nebula of ETHOS~1 adding further evidence towards the long suspected relationship between binary CSPN and jets. Their tips present emission line ratios that are consistent with shocked models of FLIERS in PNe. The fact that the jets are detached suggests a limited period of jet activity consistent with a transient accretion disk before the CE is ejected. The kinematic age of the jets ($1750\pm250$ yrs/kpc) was found to be older than the inner nebula ($900\pm100$ yrs/kpc) supporting this hypothesis. ETHOS~1 therefore continues to follow this trend previously identified in A~63 (Mitchell et al. 2007) and the Necklace (Corradi et al. 2010). Both ETHOS~1 and the Necklace have younger kinematic ages than A~63 consistent with the apparent ongoing cooling in the younger jets via [O~III] emission. A close binary engine powering jets is out of place in the literature with most models of jets incorporating orbital periods of several years (Cliffe et al. 1995; Haro-Corzo et al. 2009; Raga et al. 2009). These models may be adequate if the accretion disk is established before the companion begins its in-spiral phase as our kinematic ages suggest. Nevertheless, new models with shorter orbital periods may be necessary to model the jets in objects like ETHOS~1 which are relatively more collimated than extant model jets. VLT FORS imaging and MES high-resolution spectroscopy of the inner nebula of ETHOS~1 was conducted. The $X$-shaped inner nebula was reconstructed using a \textsc{shape} kinematic model consisting of bipolar outflows and a spherical nebula. The faint bipolar extensions particularly visible in the [O~III] FORS image could be an earlier ejection occurring at about the same time as the jets or alternatively may represent a bubble produced by the jets clearing out a cavity around the PN. Similar bipolar outflows are seen in A 65 (Walsh \& Walton 1996), PPA 1759$-$2834 (Miszalski et al. 2009b) and Fg 1 (Boffin \& Miszalski in prep.) suggesting they may have a binary-related origin. It may also be possible that ETHOS~1 represents a slightly more evolved state of Hb 12 (Vaytet et al. 2009) where the spherical nebula is ejected on top of a pre-existing Hb 12 system.

1012.3331

1012

1012.3936_arXiv.txt

{Recent observations of the low-mass pre-main sequence (PMS), eccentric spectroscopic binaries DQ\,Tau and V773\,Tau\,A reveal that their millimeter spectrum is occasionally dominated by flares from non-thermal emission processes. The transient activity is believed to be synchrotron in nature, resulting from powerful magnetic reconnection events when the separate magnetic structures of the binary components are briefly capable of interacting and forced to reorganize, typically near periastron.} {We conducted the first systematic study of the millimeter variability toward a sample of 12 PMS spectroscopic binaries with the aim to characterize the proliferation of flares amongst sources likely to experience similar interbinary reconnection events. The source sample consists entirely of short-period, close-separation binaries that possess either a high orbital eccentricity ($e\,{>}\,$0.1) or a circular orbit ($e\,{\approx}\,$0).} {Using the MAMBO2 array on the IRAM 30m telescope, we carried out continuous monitoring at 1.25\,mm (240\,GHz) over a 4-night period during which all of the high-eccentricity binaries approached periastron. We also obtained simultaneous optical VRI measurements, since a strong link is often observed between stellar reconnection events (traced via X-rays) and optical brightenings.} {UZ\,Tau\,E is the only source to be detected at millimeter wavelengths, and it exhibited significant variation ($F_{\rm 1.25mm}$\,=\,87--179\,mJy); it is also the only source to undergo strong simultaneous optical variability ($\Delta{}R\,{\approx}\,$0.9\,mag). The binary possesses the largest orbital eccentricity in the current sample, a predicted factor in star-star magnetic interaction events. With orbital parameters and variable accretion activity similar to DQ\,Tau, the millimeter behavior of UZ\,Tau\,E draws many parallels to the DQ\,Tau model for colliding magnetospheres. However, on the basis of our observations alone, we cannot determine whether the variability is repetitive, or if it could also be due to variable free-free emission in an ionized wind.} {UZ\,Tau\,E brings the number of known millimeter-varying PMS sources to 3 out of a total of 14 monitored binaries now in the literature. Important factors in the non-detection of the rest of our targets are the coarse time-sampling and limited millimeter sensitivity of our survey. We recommend that future studies concentrate on close-by targets, and obtain millimeter and optical data points with better temporal resolution. }

Recent observations of young stellar objects (YSOs) are challenging the long-standing notion that the millimeter continuum emission characterizing these objects is always dominated by the quiescent thermal emission from passively heated circumstellar dust. Powerful, transient millimeter flares attributed to synchrotron continua have now been reported toward two embedded protostars in the Corona Australis cloud \citep{choi2009}, one protostar in the Orion BN/KL star-forming region \citep{forbrich2008}, the embedded YSO GMR-A in Orion \citep{bower2003,furuya2003}, the classical T~Tauri star (CTTS) DQ\,Tau in Taurus \citep{salter2008,salter2010}, and the weak-line T~Tauri star (WTTS) V773\,Tau\,A, also in Taurus \citep{massi2002,massi2006,massi2008}. These millimeter flares are thought to be more powerful examples of the prevalent, lower-energy radio activity observed toward YSOs, and are not unlike millimeter flares occurring on the Sun \citep{stine1988,white1992}. The emission is attributed to a combination of gyrosynchrotron and synchrotron radiation powered by magnetic reconnection events in the stellar coronae, which occur when oppositely directed magnetic field lines interact \citep{bastian1998}. The radio flare resulting from this magnetic activity is expected to be accompanied by an X-ray flare of proportional luminosity according to the Neupert effect \citep{neupert1968,gudel2002}. In this way, large X-ray surveys are contributing to the characterization of the magnetic activity during the T\,Tauri phase \citep[e.g.][]{getman2005,gudel2007}, which represents the period during the formation and main sequence life of a solar-type star when magnetic activity levels are highest and when reconnection---and not accretion---is believed to be the primary X-ray production mechanism \citep{preibisch2005,stassun2006,stassun2007,forbrich2007,feigelson2007}. The X-ray data confirm coronal activity analogous to that of the Sun, but with luminosities 10$^3$--10$^5$ times higher \citep{testa2010}. Uninterrupted, long-duration observing campaigns of star-forming regions also document a once-a-week statistical occurrence of giant X-ray flares, representing the most powerful events and an estimated 1\% of all flares \citep{favata2005,getman2008a,getman2008b}. If magnetic activity is indeed the trigger, then these X-ray events are most likely accompanied by radio events due to synchrotron emission processes extending into the millimeter regime. It is noteworthy that in the two best-studied millimeter flare cases, both DQ\,Tau and V773\,Tau\,A are close-separation, eccentric, pre-main sequence (PMS) binaries with similar orbital characteristics. In addition, their flaring was recurring, leading the authors in the latest study to conclude that DQ\,Tau could display excess millimeter flux as much as 8\% of the time, with consistency near periastron. One can thus speculate that in close binaries, in addition to the single-star coronal activity described until now, two additional magnetic activity scenarios might exist. An example of the first is V773\,Tau\,A, where interbinary interactions due to chance alignments of narrow extended coronal features, like helmet streamers, cause the flares. The second binary-specific phenomenon is the current model for DQ\,Tau where colliding dipolar magnetospheres represent a simple geometric scenario for periodic events, with two primary determining factors: a periastron approach smaller than twice the magnetospheric radius ($R$\,$\sim$\,5\,$R_\star$) and a high eccentricity. In this arrangement, the closed stellar magnetospheres must overlap near periastron, but only temporarily. Once again, X-ray studies can provide constraints on the coronal extent of PMS stars, with the result that typical inferred loop lengths are 4--20\,$R_\star$ for the most powerful outbursts. This is much larger than any coronal structure observed toward more evolved stars \citep{favata2005}, and consistent with the T\,Tauri stage being the most magnetically active phase of star formation. In the case of the WTTS binary system V773\,Tau\,A, two separate coronal structures extending to $\geq$\,15\,$R_\star$ each are necessary to bridge the interbinary gap \citep{massi2008}. Toward DQ\,Tau the derived loop lengths from both X-ray and millimeter analyses are 5\,$R_\star$ in height (Getman et al.~2010, submitted; Salter et al.~2010). If the giant X-ray flare statistics correctly predict a common once-a-week occurrence with consistently large loop lengths, then more interbinary collisions might be expected in a number of close-separation binaries, in addition to any single-star events that may occur. Binary systems also occur frequently, representing 65\% or more of the local field population in the middle of the main sequence \citep{duquennoy1991}. The fraction increases to up to 75\% for the population in the Taurus star-forming region \citep{leinert1993,ghez1997,kohler1998,luhman2010}, suggesting that more systems could start out as binaries. Thus, candidate millimeter-variable systems are worthy of investigation. To assess the proliferation of significant millimeter variability among PMS binaries, we report here on a targeted millimeter variability survey of 12 PMS spectroscopic binaries that are most likely to experience millimeter flares based on predictions by the current interbinary magnetic reconnection models, either following the V773\,Tau\,A scenario and exhibiting strong flares at many orbital phases, or exhibiting the DQ\,Tau phenomena showing flares with more regularity around periastron. Since in both the DQ\,Tau and UZ\,Tau\,E cases, optical brightenings are common near periastron due to periodic accretion events \citep{jensen2007}, and because the optical light curve of DQ\,Tau was found to mirror its millimeter flare activity in both time and duration \citep{salter2010}, we complemented our millimeter data with simultaneous optical monitoring of our targets.

\label{summary} Using the IRAM 30m telescope, we have conducted a monitoring program covering 4 consecutive nights to study the millimeter variability toward 12 PMS spectroscopic binaries mostly in the Taurus and Orion star-forming regions. Here we report that one source, the CTTS UZ\,Tau\,E, experiences significant millimeter flux variations ($F_{\rm 1.25\,mm}$ ranges from 87 to 179\,mJy) on daily timescales, a clear indication of non-thermal emission processes near periastron. The rest of the sample, consisting mainly of WTTS up to three times more distant, remain undetected in the continuum for the duration of the campaign, defining upper flux limits of 5--10\,mJy at 1.25\,mm (240\,GHz). The motivation for our survey follows the recent discoveries of recurring, bright (up to 27 times quiescent values or peaking at about 0.5\,Jy) millimeter outbursts toward the T~Tauri binaries V773\,Tau\,A \citep{massi2008} and DQ\,Tau \citep{salter2010}. Attributed to synchrotron activity from interbinary interactions of large magnetic structures, the phenomenon toward V773\,Tau\,A is described as chance collisions between extended coronal features, whereas it has been proposed that the geometry of the DQ\,Tau system alone (specifically a large $e$ and small $d_{\rm min}$) results in global interactions between the two closed stellar magnetospheres near periastron \citep{mathieu1997,basri1997}. Therefore, our target list consisted of 6 close-separation binaries with circular orbits ($e$\,$\approx$\,0) that may experience activity at any time (but apparently did not do so during our observing run), as well as 6 geometrically favorable high-eccentricity systems with activity most likely to occur near periastron (where we detect two possible events toward the source UZ\,Tau\,E). In our sample, no system geometry is quite as extreme as DQ\,Tau, although our detected source UZ\,Tau\,E comes closest. Therefore, a positive detection of (double-peaked) variability near periastron toward UZ\,Tau\,E lends strong support to a similar global interbinary interaction; but does not confirm it. Instead, this detection brings our total millimeter-variable source statistics to 3 (i.e.~V773\,Tau\,A, DQ\,Tau, and UZ\,Tau\,E) out of 14 observed sources, and means that we may need to consider that millimeter flares are not so uncommon. In addition, as we examine the other systems in much greater detail, it becomes clear why we might not have expected to see, or might have missed, evidence of flares in these systems. The most important factor seems to be the flux-distance inverse relation, which affects our detection limits and sampling coverage. The study itself was also limited by the number of close-separation binaries that have been both identified and well characterized. UZ\,Tau\,E should certainly be considered for follow-up observations to help characterize the light curve profile, also on orbital timescales, and to assess potential contributions from strongly varying free-free emission processes. We must also strongly caution against the reliability of any disk model for UZ\,Tau\,E that is based on continuum flux measurements until the true quiescent flux level can be established. In the future, ALMA will allow better monitoring of these systems, leading to a more complete analysis of the proliferation of strong millimeter activity in low-mass PMS binary systems, as well as how much energy can be released during a millimeter outburst and the magnetic field regeneration timescales possible in systems known to experience recurring outbursts.

1012.3936

1012

1012.2047_arXiv.txt

We present the chemical abundance analysis of 33 red giant stars belonging to the complex stellar system Terzan~5. We confirm the discovery of two stellar populations (Ferraro et al. 2009, Nature, 462,483) with distinct iron abundances: a relatively metal-poor component with [Fe/H]$= -0.25\pm0.07$ r.m.s., and another component with [Fe/H]$=+0.27\pm0.04$ r.m.s., {\it exceeding in metallicity any known Galactic globular cluster}. The two populations also show different [$\alpha$/Fe] abundance ratios. The metal-poor component has an average [$\alpha$/Fe]$= +0.34\pm 0.06$ r.m.s., consistent with the canonical scenario for rapid enrichment by core collapse supernovae (SNe). The metal-rich component has [$\alpha$/Fe]$ = +0.03 \pm 0.04$ r.m.s., suggesting that the gas from which it formed was polluted by both type II and type Ia SNe on a longer timescale. Neither of the two populations shows evidence of the [Al/Fe] over [O/Fe] anti-correlation, that is typically observed in Galactic globular clusters. Because these chemical abundance patterns are unique, we propose that Terzan 5 is not a true globular cluster, but a stellar system with a much more complex history of star formation and chemical enrichment.

\label{intro} Terzan~5 is commonly catalogued as a globular cluster (GC) located in the inner bulge of our Galaxy. It is heavily reddened, with an average color excess $E(B-V)=2.38$ \citep{bar98,val07} and such a reddening strongly depends on the line of sight \citep{ort96,val07}. This stellar system also harbors an exceptionally large population of millisecond pulsars (MSPs): indeed, the 34 MSPs detected so-far in Terzan~5 amount to $\sim 25\%$ of the entire sample of known MSPs in Galactic GCs \citep[][ see the updated list at {\tt http://www.naic.edu/$\sim$pfreire/GCpsr.html}]{ransom05}. Recently, a combined photometric and spectroscopic study of Terzan~5 has led to the discovery of two distinct populations, as traced by two well separated ($\delta K\simeq 0.3$) Red Clumps in the $(K,J-K)$ Color Magnitude Diagram (CMD), with a $\approx$0.5 dex difference in their iron content \citep[][ hereafter F09]{fer09}. A conventional isochrone fit is consistent with the two populations of Terzan 5 being separated by a few Gyr (F09), although only a small age gap is needed if the younger population is enhanced in helium \citep{danto10}. The findings in F09 appear to be best understood if Terzan~5 was much more massive in the past than today, in order to retain the SN ejecta and igniting other star formation episodes. A more massive proto-Terzan~5 would also naturally explain its large population of MSPs and the fact that the metal-rich component is more centrally concentrated than the metal-poor one \citep[F09; see also][]{lan10}, a typical feature of stellar systems which are self-enriched in iron, as, e.g., the dwarf galaxies. With the aim of accurately reconstructing the puzzle of the formation and evolutionary history of Terzan~5, we are currently undertaking a global study of the photometric, chemical and kinematic properties of its stellar populations. This Letter presents the results of the spectroscopic screening of a suitable sample of giant stars, in order to obtain chemical abundances and abundance patterns of key metals, like iron, carbon, aluminum, oxygen and other $\alpha$-elements, and constrain the complex chemical enrichment history of Terzan~5.

The main observational evidences from the photometric and spectroscopic studies performed so far on Terzan 5 can be summarized as follows. \begin{itemize} \item Terzan~5 shows at least two stellar populations (as traced by both Red Clump and RGB stars) with distinct iron content and $[\alpha$/Fe] abundance patterns. The metal-poor population ([Fe/H]$\simeq$-0.2) is $\alpha$-enhanced and closely resembles the bulk of the old bulge population (except for the extremely small spread in iron), which formed early and quickly from a gas mainly polluted by type II SNe. The metal-rich population has a metallicity ([Fe/H]$\simeq$+0.3), and an approximately scaled Solar [$\alpha$/Fe] ratio, requiring a progenitor gas further polluted by both type II and type Ia SNe on a longer timescale. It is difficult to place the chemistry of Terzan 5 within the framework of known GCs. Indeed, while no genuine Galactic GC displays such a large difference in the iron content, and even remotely resembles the metallicity regime of the two stellar populations of Terzan~5, stars with similar iron content have been observed in the bulge field \citep[see][]{rich07,ful07,zoc08}. \item Neither Terzan~5 as a whole, nor the two populations separately show evidence of the Al-O anti-correlation. As soon as the anti-correlation is also effective at Solar metallicity and above, this further suggests that Terzan~5 as a whole is not a genuine GC, and also that it cannot be the merging of two globulars. \item Its current mass of a few $10^6 M_{\odot}$ \citep{lan10} is not sufficient to retain the SN ejecta and the large population of neutron stars which, thanks to an exceptionally high stellar collision rate \citep{verhut87,lan10}, could have been recycled into the multitude of MSPs that we observe today. \end{itemize} In order to draw more firm conclusions about the origin of Terzan~5 and its possible bimodal nature it is necessary to {\it (i)} complete the chemical screening of its populations, by also sampling stars that, in the CMD, are located between the two main RGBs, {\it (ii)} perform and analyze ultra-deep IR imaging to accurately measure the luminosity of the main sequence turn-off point(s) and derive the ages of each component, and {\it (iii)} combine radial velocity and proper motion measurements to properly determine the kinematics of the system. However, considering the information available so far, we venture the following speculations. The complex stellar population of Terzan~5 and the higher central concentration of the most metal-rich component \citep[see F09 and][]{lan10} could be naturally explained within a self-enrichment scenario. An originally more massive proto-Terzan experienced the explosions of a large number of type II and type Ia SNe, whose ejecta have been retained within the potential well and which could also have wiped out the anti-correlation signatures typical of GCs. In such a scenario, the Terzan~5 evolution should have been characterized by two main and relatively short episodes of star formation, thus accounting for the small metallicity spread of both populations. In addition, the striking chemical similarity between Terzan~5 and the bulge population can also suggest a strong evolutionary link between these two stellar systems and possibly a common origin and evolution. The current view \citep{kor04,imm04,she10} for the formation of a bulge structure suggests a range of physical processes that can be grouped in two main scenarios: (1) rapid formation occurring at early epochs (as a fast dissipative collapse, mergers of proto-clouds/sub-structures, evaporation of a proto-disk, etc.), generating a spheroidal bulge populated by old stars, and (2) evolution of a central disk/bar and its possible interaction with other sub-structures on a longer timescale. Within this framework, Terzan 5 might well be the relic of a larger sub-structure that lost most of its stars, probably because of strong dynamical interactions with other similar systems at the early epoch of the Galaxy formation, and/or later on with the central disk/bar. While most of the early fragments dissolved/merged together to form the bulge, for some (still unclear) reasons Terzan 5 survived the total disruption. Note that within this scenario, while the oldest population of Terzan 5 would trace the early stages of the bulge formation, the younger one could contain crucial information on its more recent chemical and dynamical evolution. The metal rich sub-component of Terzan 5 stands as a remarkable stellar population, worthy of more study.

1012.2047

1012

1012.2271_arXiv.txt

We present experimental results on the formation of supersonic, radiatively cooled jets driven by pressure due to the toroidal magnetic field generated by the 1.5 MA, 250 ns current from the MAGPIE generator. The morphology of the jet produced in the experiments is relevant to astrophysical jet scenarios in which a jet on the axis of a magnetic cavity is collimated by a toroidal magnetic field as it expands into the ambient medium. The jets in the experiments have similar Mach number, plasma beta and cooling parameter to those in protostellar jets. Additionally the Reynolds, magnetic Reynolds and Peclet numbers are much larger than unity, allowing the experiments to be scaled to astrophysical flows. The experimental configuration allows for the generation of episodic magnetic cavities, suggesting that periodic fluctuations near the source may be responsible for some of the variability observed in astrophysical jets. Preliminary measurements of kinetic, magnetic and Poynting energy of the jets in our experiments are presented and discussed, together with estimates of their temperature and trapped toroidal magnetic field.

Collimated outflows (jets) are associated with widely diverse astrophysical environments but exhibit many common features which are independent of the central source (\cite{livio02nature}). In general, it is believed that the ejection of jets relies on the conversion of gravitational energy into Poynting flux which powers the outflows. The standard magneto-hydrodynamic (MHD) models of jet formation rely on differential rotation along a large scale poloidal magnetic field $B_{P}$ to generate a toroidal magnetic field component $B_{\phi}$ which accelerates and collimates a disk-wind (\cite{blandford82}). Our experiments are designed to model the acceleration and collimation of astrophysical jets taking place under the condition $|B_{\phi}|$$\gg$$|B_{P}|$. Astrophysical jets and outflows are described to a first approximation by ideal MHD, which is a valid approximation when the dimensionless Reynolds ($Re$), magnetic Reynolds ($Re_M$), and Peclet ($Pe$) numbers are much larger than unity (\cite{ryutov99, ryutov00}). In this regime the transport of momentum, magnetic fields, and thermal energy, respectively, occur predominantly through advection with the flow. In this paper we wish to extend the study of supersonic (Mach number$>$1), radiatively cooled, magnetically driven (plasma beta, the ratio of the thermal to magnetic pressure, $\beta$$\lesssim$1) plasma jets from laboratory experiments which are relevant to the launching mechanism in astrophysical jet models (\cite{suzukividal09, ciardi09, lebedev05mnras}). We present new experimental results from radial foils which include the energy balance inside a magnetic cavity, temperature of the jet and trapped magnetic field inside the outflows.

We have presented new experimental data from radial foil experiments which complement previous results and enhance our understanding on the dynamics of episodic magnetically driven jets. Preliminary measurements of the kinetic and magnetic energy inside the magnetic cavities are in good agreement with the total electromagnetic energy injected as Poynting flux, measured with an inductive probe connected at the base of the cavity. The temperature of the jet/outflow, determined by time-integrated spectroscopy, resulted in $T$$\sim$170-600 eV, providing a first estimate of this parameter. These high temperatures are consistent with measurements of toroidal magnetic field inside the cavity, indicating that magnetic flux is conserved as the cavity expands and that the plasma has a magnetic Reynolds number larger than unity. Our results indicate values of the order of $Re_M$$\sim$300-1000, consistent with previous results from computer simulations (\cite{ciardi09}) thus validating the scaling of these experiments to astrophysical jets. In addition we find that episodic jet formation is halted by a metal needle on the axis of the foil. We observed the formation of only a single magnetic cavity, as in our previous experiments with radial wire arrays. This supports the hypothesis that instabilities in the jet may disrupt the current path inside the cavity, encouraging the reconnection of current at the base. Future experiments will look into the detailed dynamics of this new setup.

1012.2271

1012

1012.4322.txt

We use the secondary infall model described in Del Popolo (2009), %the results of Del Popolo (2009), using a secondary infall model which takes into account the effect of dynamical friction, ordered and random angular momentum, baryons adiabatic contraction and dark matter baryons interplay, to study how %the density profiles and inner slopes of relaxed $\Lambda$CDM dark matter (DM) halos with and without baryons (baryons+DM, and pure DM) %$\Lambda$CDM dark matter halos and baryons depend on redshift and on halo mass. We apply the quoted method to structures on galactic scales and clusters of galaxies scales. We find that the inner logarithmic density slope, $\alpha\equiv d\log\rho/d\log r$, of dark matter halos with baryons has a significant dependence on halo mass and redshift with slopes ranging from $\alpha \simeq 0$ for dwarf galaxies to $\alpha \simeq 0.4$ for objects of $M \simeq 10^{13} M_{\odot}$ and $\alpha \simeq 0.94$ for $M \simeq 10^{15} M_{\odot}$ clusters of galaxies. Structures slopes increase with increasing redshift and this trend reduces going from galaxies to clusters. %with increasing values of the mass?????? In the case of density profiles constituted just of dark matter the mass and redshift dependence of slope is very slight. In this last case, we used the Merrit et al. (2006) analysis who compared $N$-body density profiles with various parametric models finding systematic variation in profile shape with halo mass. This last analysis suggests that the galaxy-sized halos obtained with our model have a different shape parameter, i.e. a different mass distribution, than the cluster-sized halos, obtained with the same model. The results of the present paper argue against universality of density profiles constituted by dark matter and baryons and confirm claims of a systematic variation in profile shape with halo mass, for dark matter halos.

A fundamental idea that has come out of the numerical approach is that relaxed halos are (nearly) universal in many respects: %As $N$-body simulations of cold dark matter halos have become more detailed, %Several ''universal'' properties have emerged such as the phase-space density profile (Taylor \& Navarro 2001), the linear relation between the density slope and the velocity anisotropy (Hansen \& Moore 2006), the density profile (Navarro, Frenk \& White 1996 (hereafter NFW)), the distribution of axial ratio, the distribution of spin parameter, and the distribution of internal specific angular momentum (Wang \& White 2008). For what concerns the density profiles, numerous cosmological studies (e.g., Dubinski \& Carlberg 1991; Crone at al. 1994; Navarro et al. 1996, 1997; Carlberg et al. 1997; Moore et al. 1998, 1999; Jing \& Suto 2000) %citep[e.g.][]{1991ApJ...378..496D, 1994ApJ...434..402C, %1996ApJ...462..563N, 1997ApJ...477L...9F, 1997ApJ...490..493N, %1997ApJ...485L..13C, 1998ApJ...499L...5M, 1999MNRAS.310.1147M, %2000ApJ...529L..69J} revealed similar density profiles over several orders of magnitude in halo mass, with a central cusp and a $\rho(r) \propto r^{-3}$ fall--off at large radii. %However, the simulations have not been able to reach agreement about the exact behavior of the profiles in the innermost regions, where the limited %force resolution of simulations sets a lower limit to the radial range that can be probed. Various authors claim that the logarithmic slope of the %density profile reaches a value between $-1$ and $-1.5$, perhaps dependent on mass or merger history, and there is also discussion whether the %inner slope is actually universal or not %\citep{1998ApJ...499L...5M,2001ApJ...554..903K,2004MNRAS.349.1039N,2004ApJ...606..625F,2006AJ....132.2685M,2006AJ....132.2701G,2008MNRAS.387..536G}. The universal shape of the profiles, found by Navarro et al. (1996, 1997), has been confirmed by several other studies, even if the actual value of the inner density slope $\alpha$ has been a matter of controversy (Moore et al. 1998; Jing \& Suto 2000; Ghigna et al. 2000; Fukushige \& Makino 2001) and more recently, the functional form %shape of the universal profile has been substituted by profiles whose logarithmic slope becomes increasingly shallower inwards (Power et al. 2003; Hayashi et al. 2003 and Fukushige et al. 2004; Navarro et al. 2004; Stadel et al. 2008). Moreover, there is also discussion whether the inner slope is actually universal or not (Moore et al. 1998; Jing \& Suto 2000; Subramanian et al. 2000; Klypin et al. 2001; Ricotti 2003, Ricotti \& Wilkinson 2004; Cen et al. 2004; Navarro et al. 2004; Fukushige et al. 2004; Merrit et al. 2005; Merrit et al. 2006; Graham et al. 2006; Ricotti et al. 2007; Gao et al. 2008; Host \& Hansen 2009). %citep{1998ApJ...499L...5M,2001ApJ...554..903K,2004MNRAS.349.1039N,2004ApJ...606..625F,2006AJ....132.2685M,2006AJ....132.2701G,2008MNRAS.387..536G}. %Another important point on which our results can tell something is the debate on the universality of the density profiles of dark matter haloes. %As previously reported, the result that the density profiles of haloes in CDM and other hierarchical clustering %cosmologies have a universal form which is well represented %by the simple fitting formula given by NFW96, 97, has been confirmed in almost all the subsequent papers dealing with %the subject except some of them All of the previously quoted analyses dealing with universality of dark matter profiles do not study the possible effects produced by the presence of baryons, whose effect is to shallow (El-Zant et al. 2001, 2004; Romano-Diaz et al. 2008) and to steepen (Blumenthal et al. 1986; Gnedin et al. 2004; Klypin et al. 2002) the dark matter profile. In real galaxies, the two quoted effects combine, with the result of giving rise to density profiles which are different from those predicted in N-body simulations (Del Popolo 2009). In collisionless N-body simulations, this complicated interplay between baryons and dark matter is not taken into account, because it is very hard to include the effects of baryons in the simulations. However, in order to have a clear view of what simulations can tell about the inner parts of the density profiles, it is necessary to run N-body simulations that repeat the mass modeling including a self-consistent treatment of the baryons and dark matter component. The question is whether or not baryon-DM interactions are universal, or depend on scale - that could be answered by simulations in the next few years. % %%The resolution has just got to the level of 0.5\% of the virial halo including gas dynamics which is almost enough to resolve the creation of a central core. %%In other terms, we have to wait % In the present paper, we deal with the problem of the universality of density profiles made of dark matter and baryons, by using the results of Del Popolo (2009). In the next sections, we use Del Popolo (2009) model to study how inner slopes in density profiles change when baryons are present. The paper is organized as follows: in Section 2, we describe the Del Popolo (2009) model. In Section 3, we discuss the results. Finally, Section 4 is devoted to conclusions. \section[]{Model} The dark matter halos in the present study are formed using the analytical method introduced in Del Popolo (2009). In this section, we shortly summarize the model that will be used to study the behavior of the inner part of dark matter haloes when baryons are present. %The method is fully described in Del Popolo (2009). In the secondary infall model (SIM) (Gunn \& Gott 1972), %spherical collapse model of Gunn \& Gott (1972), a bound mass shell of initial comoving radius $x_i$ will expand to a maximum radius $x_m$ (named apapsis or turnaround radius $x_{ta}$). As successive shells expand to their maximum radius, they acquire angular momentum and then contract on orbits determined by the angular momentum. Dissipative physics and the process of violent relaxation will eventually intervene and convert the kinetic energy of collapse into random motions (virialization). %In reality, after reaching maximum radius, a shell collapses and will start oscillating and it will contribute to the %inner shells and so even energy is not an integral of motion anymore. The effect of the in-falling outer %shells on the dynamics of a given shell can be described assuming that the collapse is ``gentle". One can assume that the %potential well near the center varies adiabatically (Gunn 1977, Filmore \& Goldreich 1984 (hereafter FG84)), which means %that a shell near the center makes many oscillations before the potential changes significantly (Gunn 1977, FG84, %Zaroubi \& Hoffman 1993). The final density profile can be obtained in terms of the density at turn-around, $\rho_{ta}(x_m)$, the collapse factor\footnote{The collapse factor is defined as the $f=x/x_m$ (see Del Popolo 2009, Appendix A)}, and the turn-around radius (Eq. A18, Del Popolo 2009), %Eq. (\ref{eq:dturn}) together with Eq. (\ref{eq:rc}) as: \begin{equation} \rho(x)=\frac{\rho_{ta}(x_m)}{f^3} \left[1+\frac{d \ln f}{d \ln x_m} \right]^{-1} \label{eq:dturnnn} \end{equation} In our model we supposed that protostructure forms around peaks of the density field, we took into account the presence of baryons, baryons adiabatic collapse, dynamical friction, and angular momentum. These quantities can be specified as we describe in the following. The density profile of a proto-halo is taken to be the profile of a peak in a density field described by the Bardeen et al. (1986) power spectrum, as is illustrated in Del Popolo (2009), Figure 6. In the present paper, we take into account the ordered angular momentum, $h$, (Ryden \& Gunn (hereafter RG87)) which arises from tidal torques experienced by proto-halos, %and is usually quantified as a dimensionless spin parameter $\lambda$ (Peebles 1969), and the random angular momentum, $j$, (RG87)) which is connected to random velocities (see RG87). %: random $j$, and ordered, $h$. Ordered angular momentum is got obtaining, first, the rms torque, $\tau (x)$, on a mass shell, then obtaining total specific angular momentum, $h(r,\nu )$, acquired during expansion by integrating the torque over time (Ryden 1988a (hereafter R88), Eq. 35) (see Appendix C of Del Popolo 2009). The random part of angular momentum was assigned to protostructures according to Avila-Reese et al. (1998) scheme. This consists in expressing the specific angular momentum $j$ through the ratio $e_0=\left( \frac{r_{min}}{r_{max}} \right)_0$, where $r_{min}$ and $r_{max}$ are the maximum and minimum penetration of the shell toward the center, respectively, %COSA SONO $r_min$, $r_max$???? and left this quantity as a free parameter (see Appendix C of Del Popolo 2009). %The effect of dynamical friction on the central halo profile has been studied in several papers (e.g., %El-Zant et al. 2001 (hereafter EZ01); Tonini, Lapi \& Salucci 2006 (hereafter TLS)) using different approaches. EZ01 %studied the changes induced by dynamical friction in the dark matter halo structure within the context of galaxy %formation, where gas cools to form dense clumps. The orbital energy lost by the clumps to the dark matter background is %sufficient to ``heat up" and flatten the dark matter density cusps. In the present paper, we took into account dynamical friction by introducing the dynamical friction force in the equation of motion (see Del Popolo 2009, Eq. A16). Dynamical friction force was calculated dividing the gravitational field into an average and a random component generated by the clumps constituting hierarchical universes following Kandrup (1980) (see Appendix D of Del Popolo 2009). The shape of the central density profile is influenced by baryonic collapse: baryons drag dark matter in the so called adiabatic contraction (AC) steepening the dark matter density slope. % %%For a spherically symmetric protostructure that consists of a fraction $F << 1$ (this fraction is fixed using WMAP %%collaboration data (Spergel et al. 2003)) of dissipational baryons and a fraction $1-F$ of dissipationless dark matter %%particles constituting the halo, Blumenthal et al. (1986) described an approximate analytical model to calculate the %%effects of AC. More recently Gnedin et al. (2004) proposed a simple modification of the Blumenthal model, which %%describes numerical results more accurately. %%For systems in which angular momentum is exchanged between baryons and dark matter (e.g., through dynamical friction), %%Klypin et al. (2002) introduced a modification to Blumenthal's model. % The adiabatic contraction was taken into account by means of Gnedin et al. (2004) model and Klypin et al. (2002) model taking also account of exchange of angular momentum between baryons and dark matter (see appendix E of Del Popolo 2009 for a wider description). Our method of halo formation has considerable flexibility with direct control over the parameter space of initial conditions differently from numerical simulations which yield little physical insight beyond empirical findings precisely because they are so rich in dynamical processes, which are hard to disentangle and interpret in terms of underlying physics. %For example, we can change the shape of the primary density peak by smoothing or altering the matter power spectrum.

In the present paper, we studied the dependence of the inner slope of the density profiles of dark matter halos with and without baryons. The density profiles were built up using the Del Popolo (2009) method, which includes the effect of baryons on dark matter halos. We calculated the inner slopes for three different cases: in the case A we take into account all the effects included in Del Popolo (2009), namely angular momentum, dynamical friction, baryons, baryons adiabatic contraction; in the case B there are no baryons and dynamical friction, and angular momentum is taken into account as in Del Popolo (2009); in case C only angular momentum is taken into account with magnitude reduced as in Del Popolo (2009) in order to reproduce the angular momentum of N-body simulations. In case A, and B, we found a strong dependence of the inner slope with mass and redshift: lower mass objects having smaller slopes which increases with redshift. This result argues against the universality of these kind of profiles. When baryons are not present (case C), the mass and redshift dependence of slope is quite slight in agreement with results of other studies (e.g., Gao et al. 2008). However, to have some further insights on the variation of the profile shape with the halo mass, we used a method similar to that of Merrit et al. (2006), who plotted their $N$-body halos, in the profile shape vs. mass plane. %%%As them, we used the profile shape parameter $n$ from the S\'ersic model fit to the light profile, or the corresponding parameter from the Prugniel-Simien model fit to %%%the dark-matter density. %QUESTI HANNO FATTO I FIT ALLE GALASSIE VERE CON SERSIC E PRUGNIEL AALLA DM!!!!!!!!!!!!!!! Similarly to Merrit et al. (2006), the galaxy-sized halos obtained with our model have a different shape parameter, i.e. a different mass distribution than the cluster-sized halos obtained with the same model. The sample of dwarf- and galaxy-sized halos studied has a mean ($\pm$ standard deviation) profile shape $n=3.72 (\pm 0.8)$, while the cluster-sized halos had $n=2.85 (\pm 0.6)$. A Student $t$ test, without assuming equal variance in the two distributions, reveals that the means in our model are different at the 99.90\% level. This leads to conclude that there is a significant mass dependence in the density profiles of the halos in our model similarly to the simulated dark-matter halos in Merrit et al. (2005, 2006).

1012.4322

1012

1012.1722_arXiv.txt

The aim of the new generation of radio synthesis arrays such as LOFAR and SKA is to achieve much higher sensitivity, resolution and frequency coverage than what is available now, \san{especially at low frequencies}. To accomplish this goal, the accuracy of the calibration techniques used is of considerable importance. Moreover, since these telescopes produce huge amounts of data, speed of convergence of calibration is a major bottleneck. The errors in calibration are due to system noise (sky and instrumental) as well as the estimation errors introduced by the calibration technique itself, which we call ``solver noise''. We define solver noise as the ``distance'' between the optimal solution \san{(the true value of the unknowns, uncorrupted by the system noise)} and the solution obtained by calibration. We present the Space Alternating Generalized Expectation Maximization (SAGE) calibration technique, which is a modification of the Expectation Maximization algorithm, and compare its performance with the traditional Least Squares calibration based on the level of solver noise introduced by each technique. For this purpose, we develop statistical methods that use the calibrated solutions to estimate the level of solver noise. The SAGE calibration algorithm yields very promising results both in terms of accuracy and speed of convergence. The comparison approaches that we adopt introduce a new framework for assessing the performance of different calibration schemes.

\label{sec:introduction} Early radio-astronomy predominantly used single-dishes for observations. With the resolution requirements increasing, the single dish approach became impractical. This paved the path for using radio-interferometric techniques with multiple antennas linked together as an array that operates as a large effective single-dish \citep{A.R.1}. The sensitivity of an interferometer is greatly increased, compared to a single-dish telescope, due to the larger combined collecting area. The currently planned or built radio interferometers, such as the \san{Square Kilometer Array (SKA)\footnote{http://www.skatelescope.org}, the Murchison Widefield Array (MWA)\footnote{http://www.mwatelescope.org}, the Precision Array to Probe Epoch of Reionization (PAPER)\footnote{http://astro.berkeley.edu/\~{}dbacker/eor}, the 21-cm Array (21CMA)\footnote{http://21cma.bao.ac.cn}, the Hydrogen Epoch of Reionization Array (HERA)\footnote{http://www.reionization.org}, the Long Wavelength Array (LWA)\footnote{http://lwa.unm.edu} and the LOw Frequency ARray (LOFAR)\footnote{http://www.lofar.org}, consist of a large number of elements and include short, intermediate and many of them longer antenna spacings}. For an introduction to radio interferometry we refer the reader to \cite{A.R.1}. In the interferometric visibilities there always exist errors introduced by the sky, the atmosphere (e.g. troposphere and ionosphere), the instrument (e.g. beam-shape, frequency response, receiver gains etc.) and by Radio Frequency Interference (RFI). The process of estimating and reducing the errors in these measurements is called ``calibration'' and is an essential step before imaging the visibilities. The classical calibration method, named external (or primary) calibration, is based on observing a celestial radio source with known properties. This approach is strongly dependent on the accuracy with which the source properties are known. The external calibration is improved by using self-calibration \citep{selfcal} which utilizes the observed data for estimating both the unknown instrumental and the sky parameters. The quality of calibration and the imaging is significantly increased by iterating between the sky and the instrument model. \san{The redundant calibration is also independent of the sky model. It calibrates for both the sky and the instrument, using redundant information in the measured data. However, its performance is limited to arrays with a regular arrangement in their antennas layout.} Calibration is an optimization process that is non-linear by nature. It is in essence a Maximum Likelihood (ML) estimation of the unknown parameters by applying non-linear optimization techniques. Traditional calibration is estimating the ML solution by the non-linear Least Squares (LS) method via various gradient-based techniques such as the Levenberg-Marquardt (LM) algorithm \citep{K.L1, A.L.1}. This approach was improved by the Expectation Maximization (EM) algorithm \citep{M.2} and later on by the Space Alternating Generalized Expectation Maximization (SAGE) technique which was introduced by \citet{J.A.1} and was applied to interferometer calibration by \citet{S.2}. The analysis and application of the aforementioned schemes for the calibration of radio interferometers can be found in \citet{S.2}. To reach the scientific goals of the new generation of radio arrays, calibration algorithms must have the highest accuracy possible. Furthermore, the number of measured visibilities that has to be calibrated is unprecedented. The speed of convergence of the calibration processes must therefore be the fastest with the minimum possible computational cost. Based on these facts, the best calibration method is referred to as the one which minimizes the ``distance'' between the true values of unknown parameters and the values obtained by calibration and minimizes computational time. We should take into account that the measured data of an interferometer is always corrupted by different sources of noise such as the thermal noise, which is an additive Gaussian random process, and confusion noise \citep{condon}, which affects the coherency matrix (see next section and e.g. \citet{bornwolf}). For a detailed discussion on the sources of noise the reader is referred to the Chapter 6 of \citet{stef}. When the calibration process of the measured data is done, the ``distance'' between the true value of the unknown parameters and their calibrated solutions depends on the initial noise and the errors originating from the calibration process itself (e.g. converging to a local minimum), which is called ``solver noise''. In other words, because the calibrated solutions are not optimal, there always exists some solver noise between these solutions and the true values of the unknown parameters affected by the initial noise. The lower the solver noise, the higher the accuracy of the calibrated results. Thus, in order to increase the calibration efficiency, we need to choose the calibration scheme which has the minimum solver noise as well as the lowest computational cost. To achieve this, we should be able to compare these two factors between various calibration techniques. We introduce a general framework for detecting the level of solver noise in calibration algorithms based only on their solutions. In this paper, we present the SAGE calibration method and emphasize its superiority, compared to the traditional LS calibration, in terms of accuracy and speed of convergence. Mathematical derivations of the algorithms are presented in the appendices. We also investigate the applicability of two well-known measures, the Kullback-Leibler Divergence (KLD) \citep{K.L.D} and the Likelihood Ratio Test (LRT) \citep{L.R.T}, in revealing the level of solver noise in calibrated solutions. Illustrative examples of both real and simulated observations, show the superior performance of the SAGE calibration compared to the LS one. \san{They also indicate that the LRT approach is very promising at detecting the level of solver noise in the obtained calibrated solutions, while the KLD approach is not always conclusive.} The following notations are used in this paper: Bold, lowercase letters refer to column vectors, e.g., {\bf y}. Upper case bold letters refer to matrices, e.g., {\bf C}. All parameters are complex numbers, unless stated otherwise. The inverse, transpose, Hermitian transpose, and conjugation of a matrix are presented by $(.)^{-1}$, $(.)^T$, $(.)^H$, and $(.)^*$, respectively. The statistical expectation operator is referred to as $E\{.\}$. The matrix Kronecker product and the proper (strict) subset are denoted by $\otimes$ and $\subsetneq$, respectively. The diagonal matrix consisting of only the diagonal entries of a square matrix is given by diag$(.)$. $\bf{I}$ is the identity matrix and $\varnothing$ is the empty set. The Kronecker delta function is presented by $\delta_{ij}$. $\mathbb{R}$ and $\mathbb{C}$ are the sets of Real and Complex numbers, respectively. The Frobenius norm is shown by $||.||$. Estimated parameters are denoted by a hat, $\widehat{(.)}$. All logarithmic calculations are to the base $e$.

\label{sec:summary} Since the new generation of radio synthesis arrays are producing a large amount of data with high sensitivity, it is of great interest to devise new calibration techniques in order to increase the accuracy of solutions with the highest possible speed of convergence. In this paper, we presented the superior performance of the SAGE calibration scheme compared with the traditional LS calibration method. The superiority is in the sense that SAGE calibration has the highest accuracy, the fastest speed of convergence, and the cheapest computational cost. Since both the algorithms are estimating the ML of unknowns in different ways, it is possible that in some special cases ,such as having a very low initial noise in the measured visibilities, we don't have a specific difference between the accuracy of their solutions. While, even in this case, the SAGE calibration's faster speed of convergence and cheaper computational cost justify its application instead of the LS calibration. We compared the accuracy and the rate of convergence of the SAGE and the LS calibration in a simulated observation example. More accurate results in a much shorter time are obtained by the SAGE algorithm compared with the LS. The challenge in improving the performance of the SAGE calibration technique is to find the best way of partitioning over the unknown parameter space. This can highly affect both the accuracy and speed of convergence of the calibration process. On the other hand, there always exists some estimation errors in the calibrated solutions. These errors are originated from the system noise (sky and instrumental) in the measurements, plus ``solver noise'' which is referred to errors produced by the calibration algorithm itself. The more accurate the calibrated solutions are, the less the amount of solver noise is. Based on this fact, the best calibration method is the one which provides us with the minimum solver noise. KLD and LRT are utilized to reveal the level of solver noise in the solutions of different calibration schemes. We showed in \san{illustrative} examples that the LRT algorithm produces a very promising result. The KLD method is rather inconclusive according to the initial assumption for the PDFs fitted to the solutions. We assumed that the distribution of the solutions is a mixture of Gaussian distributions. However, in reality, the solutions may follow a different distribution and subsequently the KLD result \san{may not have the same efficiency as the LRT's}. Therefore, initially we should find the proper distribution which is appropriate for the solutions in order to calculate the KLD. The main direction of future work should be to investigate the application of the proposed calibration technique to real data obtained by LOFAR. Since LOFAR is observing the whole sky, the number of radio sources in the sky model will be very large. Subsequently, for applying the SAGE calibration, partitioning over the unknowns by manually checking the characteristics of all the sources will not be efficient. Therefore, the first challenge in utilizing the SAGE calibration to real data would be automating the partitioning over the unknowns for any given problem. Furthermore, we saw in the paper that all the mentioned calibration algorithms involve essentially the solution of a non-linear optimization problem. Applying suitable regularization techniques using proper smoothing functions to improve the accuracy of the solutions is an issue that must be investigated further in the near future. Add to that, that all the calibration schemes have the possibility of converging to a local optimum. Utilizing probabilistic techniques such as Simulated Annealing (SA) \citep{S.5131983} to assure that we are converging to a global optimum has the problem of decreasing the speed of convergence. Providing extra constraints for the mentioned calibration schemes which can guarantee the convergence to the real solutions could be one of the challenging areas of research in the future. Moreover, we have shown that the solver noise criterion could be used for revealing the level of accuracy in the calibrated solutions. Investigating possible systematic effects on the solver noise as well as the level of their influence are amongst the main issues for the future work.

1012.1722

1012

1012.4107_arXiv.txt

The two fastest near infrared survey telescopes are UKIRT-WFCAM and VISTA. The data from both these instruments are being archived by Wide Field Astronomy Unit (WFAU) at the IfA, Edinburgh, using the same curation pipeline, with some instrument specific processing. The final catalogues from these surveys will contain many tens of billions of detections. Data are taken for a range of large surveys and smaller PI programmes. The surveys vary from shallow hemisphere surveys to ultra deep single pointings with hundreds of individual epochs, each with a wide range of scientific goals, leading to a wide range of products and database tables being created. Processing of the main surveys must allow for the inclusion of specific high-level requirements from the survey teams, but automation reduces the amount of work by archive operators allowing a higher curation efficiency. The decision making processes which drive the curation pipeline are a crucial element for efficient archiving. This paper describes the main issues involved in automating the pipeline.

The WFCAM and VISTA Science Archives \citep{Hamb08} are the main access for data from WFCAM \citep{Cas07} and VISTA \citep{Emer10}. The majority of time on both instruments is spent on large surveys: UKIDSS \citep{Lawr07}, and the VISTA Public Surveys \citep{Arna07}. There are also a range of smaller Principal Investigator (PI) programmes allocated by the Telescope Allocation Committees each semester that require curating. We run the same set of tasks on all the surveys and programmes, with the amount of processing in each task dependent on the type of programme. For instance a wide survey will spend more time on band-merging and neighbour tables, but a deep survey will spend more time on deep stack creation and multi-epoch tables. PI programmes are set up completely automatically\footnote{Occasionally we have manually grouped together several related PI programmes from different semesters before automatic processing.} because of the large number of programmes, but surveys are set up in a semi-automatic way receiving special instructions, quality control and sometimes additional products from the science teams. PI programmes usually obtain all their data in one observing semester and are processed completely at the end of the semester, so new releases will be due to software or calibration improvements necessitating a complete reprocessing. Surveys build up data over many semesters and will be appended to, as well as occasionally reprocessed. These different scenarios have to be factored into the pipeline control.

1012.4107

1012

1012.2517_arXiv.txt

Highly collimated, relativistic outflows are commonly observed in compact astronomical objects (gamma-ray bursts, active galactic nuclei, galactic black-hole and neutron-star binaries). It is widely believed that in all these cases, the jet is driven by rotating, twisted magnetic fields (see, e.g., recent reviews in ref.\cite{jets}). The rapidly spinning central body (neutron star, accretion disk, black hole) twists up the magnetic field into a toroidal component and the plasma is ejected by the magnetic tension. In relativistic flows, the energy per particle significantly exceeds the rest mass energy therefore in order to create a relativistic jet, one has to convey a significant energy to a small amount of the matter. The main advantage of the magnetic launch mechanism is that the magnetic field lines, like driving belts, could in principle transfer the rotational energy to a low density periphery of the central engine thus forming a baryon pure but energetic outflow. The relativistic velocities could potentially be achieved if the magnetic energy density in the plasma frame exceeds the plasma energy density. In such outflows, the energy is transported, at least initially, in the form of the Poynting flux. The question is how and where the electromagnetic energy is eventually transformed to the plasma energy. In the scope of ideal MHD, the energy could be transferred to the plasma only via gradual acceleration by electromagnetic stresses. In non-relativistic flows, the plasma is accelerated centrifugally when sliding along the rotating poloidal field lines. The azimuthal field is generated only when the plasma inertia becomes comparable with the magnetic stresses so that the field lines have to bend backwards. This implies that a good fraction of the Poynting flux is converted into the kinetic energy of the flow already when the azimuthal field becomes comparable with the poloidal one. In relativistic, highly magnetized flows, the magnetic force is generally balanced not by inertia but by the electric force. The azimuthal field becomes comparable with the poloidal one, both of them being comparable with the electric field, at the light cylinder defined as a surface on which the corotational velocity is equal to the speed of light (in differentially rotating magnetospheres, this surface is not a cylinder but we retain the standard term, which has come from the pulsar theory). The fluid kinetic energy remains small at the light cylinder. Beyond the light cylinder, the conservation of the magnetic flux implies that the poloidal magnetic field decreases as $1/r^2$, where $r$ is the cylindrical radius of the jet. The azimuthal field and the electric field decrease only as $1/r$ being close to each other so that the fluid is only slowly accelerated by a small residual force. By this reason, the acceleration zone is extended well beyond the light cylinder so that formation of relativistic jets spans a very large range of scales. In recent studies, both numerical and analytical \cite{komissarov07,komissarov09,tchekhovskoy08,tchekhovskoy09,tchekhovskoy10,lyubarsky09,lyubarsky10}, the general conditions were formulated for collimation and acceleration of relativistic MHD jets and the efficiency of the Poynting flux into the kinetic energy conversion was thoroughly examined. Non-steady Poynting dominated outflows have also being studied \cite{granot_etal10,levinson10,lyutikov10a,lyutikov10b}. In these works, only cold flows have been addressed. Here we relax this assumption and study relativistically hot, highly magnetized jets. By relativistically hot we mean the fluid with the pressure exceeding the rest mass energy density, $p>\rho c^2$. Such outflows are believed to be formed in gamma-ray bursts (see, e.g., reviews \cite{piran04,meszaros06}). Here we study such outflows in the far zone where most of acceleration occurs. When the flow expands, the fluid cools down. In the pure hydrodynamical case, the thermal energy is converted into the kinetic energy of the fluid. It follows immediately from the Bernoulli equation that the relativistically hot fluid with the adiabatic index $\Gamma=4/3$ is accelerated such that the flow Lorentz factor grows proportionally to the jet radius. In a magnetized flow, the thermal energy could be converted not only into the kinetic energy but also into the Poynting flux. The aim of this study is to explore the fate of the thermal energy in expanding, Poynting dominated jets. We consider outflows confined by the pressure of the external medium because only in this case the jet could be collimated. The paper is organized as follows. In section 2, the general equations governing ideal relativistic axisymmetric flows are presented. In section 3, asymptotic theory is developed describing such outflows in the far zone. In sect. 4, the theory is applied to narrow jets from rigidly rotating sources. Conclusions are presented in sect. 5.

In this paper, we studied hot, strongly magnetized, relativistic jets at large distances from the source. The acceleration zone of relativistic, Poynting dominated jets spans a very large range of scales therefore the processes far away from the source are of special interest. Multi-scale systems generally pose a strong challenge to numerical simulations. On the other hand, they are suitable for asymptotic analysis. We derived equations governing the flow in the far zone; they do not contain terms that nearly cancel each other, as the general MHD equations do, and therefore they could be solved relatively easily. We concentrated on relativistically hot flows because cold, Poynting dominated jets has already been extensively studied. Our results directly generalize the results of ref. \cite{lyubarsky09} for cold flows. Making use of the obtained asymptotic equations, we studied in detail the structure of the jet from a rigidly rotating source, $\Omega=\it const$. The collimation and acceleration of the flow are intimately connected and determined by the distribution of the confining pressure. We adopted a power law distribution, $p_{\rm ext}\propto z^{-\kappa}$, so that the solution depends on two parameters, the index $\kappa$ showing how fast the pressure decreases, and the parameter $\beta$ defined by Eq. (\ref{beta_lin}) and showing how strong is the external pressure extrapolated to the light cylinder. We have found that hot, Poynting dominated jets are collimated exactly as in the cold case. Namely, at $\kappa\leq 2$ the flow shape is described by a power law function ( Eq. (\ref{Ykappa<2}) for $\kappa<2$ and Eqs. (\ref{Ykappa=2a}) and (\ref{Ykappa=2b}) for $\kappa=2$) so that the jet opening angle continuously decreases. At $\kappa>2$ the flow becomes asymptotically radial, the final collimation angle being presented by Eq. (\ref{Theta}). When the flow expands, the fluid cools down. We have shown that at $\kappa<2$ or $\kappa=2$, $\beta>1/4$, the thermal energy is converted into the kinetic energy of the flow. On the contrary, at $\kappa=2$, $\beta<1/4$ or $\kappa>2$, the thermal energy is converted into the Poynting flux so that even the relativistically hot flow is accelerated as if it is cold. Note that acceleration regimes of cold flows are also different in these parameter ranges. At $\kappa\le 2$, the cold flows are accelerated until the equipartition between the Poynting and the plasma kinetic energy fluxes is eventually achieved. At $\kappa>2$, the flow acceleration ceases when a limiting Lorentz factor is achieved so that the flow could remain Poynting dominated \cite{lyubarsky09}. One now sees that this conclusion remains valid also for relativistically hot but Poynting dominated flows. In gamma-ray bursts, the relativistic jet is formed during the collapse of star's core. In this case, the outflow is initially relativisticallly hot. Observations of the burst afterglows suggest that the final opening angle of the jet is a few degrees whereas the final Lorentz factor is at least a few hundreds. This implies $\Theta\gamma\gg 1$, which is characteristic for the case $\kappa>2$. We have shown that in this case the thermal energy is converted into the Poynting flux, not into the kinetic energy. Therefore the thermal acceleration could not help to transform the Ploynting flux into the plasma energy. Some sort of magnetic dissipation is necessary in order to utilize the electro-magnetic energy completely. \vskip 1 cm This work was supported by the US-Israeli Binational Science Foundation under grant number 2006170 and by the Israeli Science Foundation under grant number 737/07.

1012.2517

1012

1012.0724_arXiv.txt

Gravitational instability plays an important role in driving gas accretion in massive protostellar discs. Particularly strong is the global gravitational instability, which arises when the disc mass is of order 0.1 of the mass of the central star and has a characteristic spatial scale much greater than the disc's vertical scale-height. In this paper we use three-dimensional numerical hydrodynamics to study the development of gravitational instabilities in a disc which is embedded in a dense, gaseous envelope. We find that global gravitational instabilities are the dominant mode of angular momentum transport in the disc with infall, in contrast to otherwise identical isolated discs. The accretion torques created by low-order, global modes of the gravitational instability in a disc subject to infall are larger by a factor of several than an isolated disc of the same mass. We show that this global gravitational instability is driven by the strong vertical shear at the interface between the disc and the envelope, and suggest that this process may be an important means of driving accretion on to young stars.

\label{sec:intro} Accretion discs play a fundamental role in many aspects of astrophysics. Objects as diverse as planets, stars and super-massive black holes are all thought to acquire significant fractions of their mass through disc accretion, and consequently understanding accretion disc physics is important in understanding the formation of all of these objects. Critical to our understanding of accretion discs is the process of angular momentum transport, but despite many years of research on this subject we still do not fully understand the mechanism(s) by which angular momentum is transported in gaseous discs. In many cases we believe that magnetohydrodynamic instabilities, such as the magnetorotational instability \citep[MRI,][]{bh91,bh98} is the dominant transport mechanism. However, some systems, notably protostellar discs, are insufficiently ionized for the MRI to operate everywhere \citep[e.g.,][]{g96}, and it is also not clear whether or not the MRI can drive accretion rates as high as those which are observed \citep*{kpl07}. Consequently, it is still desirable to investigate other mechanisms for angular momentum transport in discs. One such mechanism which has received considerable interest in recent years is angular momentum transport by gravitational instabilities \citep[GIs; see, e.g.,][and references therein]{d07,lodato08}. Gaseous discs in Keplerian rotation become unstable to self-gravity when the \citet{toomre64} $Q$ parameter is less than some critical value of order unity. The Toomre parameter is defined as \begin{equation} \label{eq:toomreq} Q = \frac{c_{s} \Omega}{\pi G \Sigma} \, , \end{equation} where $c_{s}$ is the sound speed of the gas, $\Omega$ is the angular frequency and $\Sigma$ is the surface density. Shearing discs generally become unstable to non-axisymmetric perturbations before axisymmetric ones, so GIs in discs initially manifest themselves as spiral density waves. It was recognised long ago that such spiral density waves can transport angular momentum \citep{lbk72} but detailed study of the non-linear development of GIs in gaseous discs has only recently become possible. This process has been studied in great detail using numerical hydrodynamics, and we now have a well-established picture whereby angular momentum transport by GIs is primarily governed by disc thermodynamics. GIs in isolated thin gas discs tend to evolve to a self-regulating state, where the energy liberated by accretion is balanced by local (radiative) cooling \citep*[e.g.,][]{g01,lr04,mejia05,clc09}, and although gravity is a long-range force, ``global'' effects generally do not dominate unless the disc mass is an appreciable fraction ($\gtrsim 25$\%) of the mass of the central object \citep*{lka98,lr05}. In this picture the efficiency of angular momentum transport can be parametrized in terms of a classical \citet{ss73} $\alpha$-prescription \citep{g01,lr04}, where \begin{equation} \alpha_{\mathrm {GI}} = \frac{4}{9}\frac{1}{\gamma (\gamma-1)t_{\mathrm {cool}}\Omega} \, . \end{equation} Here $t_{\mathrm {cool}}$ is the local cooling time-scale and $\gamma$ is the adiabatic index of the gas. Faster cooling leads to deeper spiral density waves (i.e., with higher density contrasts), and thus to more efficient transport of angular momentum. However, if the cooling becomes too rapid the disc is unable to maintain its self-regulating state, and the GIs instead lead to fragmentation of the disc \citep{g01,r03}. This in turn imposes a maximum efficiency at which angular momentum can be transported by GIs without leading to disc fragmenting, and numerical simulations place typically this ``fragmentation boundary'' at $\alpha_{\mathrm {GI}} \lesssim 0.1$ \citep*[corresponding to $\beta = t_{\mathrm {cool}}\Omega \gtrsim 3$--5, e.g.,][]{g01,r03,rla05}. When extended to consider discs with realistic opacities, these results imply a maximum accretion rate that can be sustained by GIs in a self-regulating state \citep[e.g.,][see also Fig.~\ref{fig:mdot_max}]{levin03,ml05,levin07,clarke09,rafikov09}. Except at very small radii this rate is low, $\sim 10^{-6}$M$_{\odot}$yr$^{-1}$, and this raises questions as to how many astrophysical objects are able to accrete their mass in a plausible time-scale. The star may continue accreting bound clumps of gas even after the disc fragments \citep[eg.,][]{vb10}, but the details of this process remain uncertain. \begin{figure} \centering \resizebox{\hsize}{!}{ \includegraphics[angle=270]{fig1.ps} } \caption{Maximum sustainable accretion rate in a critically self-gravitating disc with $\alpha = 0.1$, computed following the procedure described in \citet{levin07} and using the opacities $\kappa(\rho,T)$ of \citet{bl94} and \citet{bell97}. The sharp jump in the critical accretion rate at an orbital period of $\simeq 300$yr is caused by the transition between the optically thick inner disc and optically thin outer disc \citep{ml05}. This corresponds to a radius of $\simeq 40$AU for a 1M$_{\odot}$ central star, or $\simeq 100$AU for a 10M$_{\odot}$ star. The maximum sustainable accretion rate at larger radii is small, $\sim 10^{-6}$M$_{\odot}$yr$^{-1}$; in the ``local limit'', larger accretion rates lead to fragmentation. In some cases external irradiation can be the dominant source of heating, imposing a temperature ``floor'' (denoted by $T_{\mathrm {min}}$) and enhancing the maximum accretion rate. Similar figures can be found in \citet{clarke09} and \citet{rafikov09}.} \label{fig:mdot_max} \end{figure} To date most numerical studies of GIs have looked at isolated self-gravitating discs, but in reality it seems likely that most gravitationally unstable discs will still be subject to some level of infall on to the disc. Indeed, in many cases it is likely that the instantaneous infall rate on to the disc exceeds the accretion rate through the disc. For example, observed accretion rates on to protostellar discs are typically an order of magnitude larger than the accretion rates on to the protostars themselves \citep*[e.g.,][]{kenyon90,calvet_ppiv}. Similar discrepancies between infall and disc accretion rates have been found in models of low-mass star formation \citep[e.g.,][]{vorobyov09}, and in models of star formation in black hole accretion discs \citep[e.g.,][]{ml04}. In this paper we present an initial investigation of this problem, by using three-dimensional numerical hydrodynamics to follow the evolution of a self-gravitating accretion disc subject to quasi-spherical infall. In Section \ref{sec:method} we present our numerical method, and in Section \ref{sec:results} we discuss the results of our simulations. We find that infall on to the disc can substantially enhance the efficiency of angular momentum transport, through the excitation of low-order, global, spiral density waves. We discuss the consequences of this result for real astrophysical systems, along with the limitations of our analysis, in Section \ref{sec:dis}, and summarize our conclusions in Section \ref{sec:conc}.

\label{sec:dis} \subsection{Limitations} We have presented calculations on the effect of infall on to a gravitationally unstable accretion disc, but our approach is highly idealised. This approach allows us to study the key physical processes in detail, but imposes some limitations when we apply our results to real astrophysical situations. Our first major simplification is in the initial conditions of our infall model. In order to ensure that we understand the various numerical effects in our calculations, we chose to let a rotating cloud fall on to an already-present disc, rather than letting the disc form self-consistently. The advantages of our set-up are two-fold. First, we can control the accretion rate on to the disc, and ensure that the ratio between infall rate and theoretical accretion rate is approximately constant over the duration of our simulation (as seen in Fig.\ref{fig:infallrate}). Second, we can compare our results directly to our isolated disc models, where the physics is well-understood, and thus isolate the effects of infall from the myriad of other potential effects. However, the trade-off is that the simulations are highly idealised, and not always realistic. In particular, we note that quasi-spherical infall is only expected in the early stages of protostellar collapse, when the disc and envelope masses are likely to be much larger than those considered here \citep[e.g.,][]{boley09,vorobyov09}. Our results have relevance to almost any case of infall on to a gravitationally unstable disc (as any infall is, by definition, sub-Keplerian), but we note that care be taken when applying our results to real systems. Our second major simplification lies in our treatment of the disc thermodynamics. Our scale-free cooling law has previously been studied in great detail \citep{g01,r03,lr04,clc09}, but it is recognised to be a poor approximation to real systems. The effect of our scale-free cooling law can be seen in Fig.~\ref{fig:panels_radprofile}: the disc temperature increases slightly with radius. Real discs almost invariably have cooling time-scales that are shorter (relative to the local dynamical time-scale) at large radii than at small radii, and are thus not scale-free \citep[e.g.,][]{rafikov07,clarke09}. A more realistic treatment would require an opacity-based cooling prescription \citep[e.g.,][]{boley06,s07}, but this would introduce several new free parameters to the problem. Again, our simplifications are not entirely physical in this regard, but do allow us to study the important processes in detail. Essentially we have chosen to perform a well-controlled numerical experiment instead of a physically realistic simulation, and our results should be interpreted with this in mind. \subsection{Comparison to Previous Work} The majority of previous work in this area has studied the transport properties of isolated gravitationally unstable discs \citep{lb94,lr04,rla05,boley06,clc09}. Our reference simulations exhibit the same behaviour as in these previous studies, as discussed in Section \ref{sec:disc_only}. However, the influence of infall on the evolution of GIs has not yet been explored in great detail. \citet{krumholz07} studied angular momentum transport in a self-consistently-formed disc subject to a very high rate of infall. They found very high accretion rates, equivalent to $\sim$30\% of the total disc mass per dynamical time-scale, with effective $\alpha$-values that exceeded unity. Most of the power was found in the $m=1$ mode, and \citet{krumholz07} attributed this very rapid accretion to the SLING instability \citep{adams89,shu90}. We note, however, that the discs formed in these simulations were much more massive than those considered here, with $q \sim 0.5$--1. It has long been known that low-order spiral modes can drive rapid accretion in massive discs \citep[e.g.,][]{lka98}, and in this regards the results of \citet{krumholz07} are not directly comparable to those of our simulations. In addition, a number of recent studies have used one- and two-dimensional simulations to study the formation and evolution of protostellar discs \citep[e.g.,][]{hueso05,vb07,vb09,vorobyov09,vd10,zhu09,zhu10}. These simulations are less computationally intensive than 3-D simulations, and are therefore able to follow the evolution of the system for much longer time-scales. They also make use of more physically realistic prescriptions for both infall and thermodynamics, forming discs self-consistently from collapsing clouds and incorporating realistic models for radiative heating and cooling. These simulations generally predict that most of the GIs' power is found in low-order spiral modes, due to both infall and the relatively high masses of the discs which form. In addition, many of these simulations have been seen to exhibit transient accretion outbursts, triggered in some cases by the accretion of bound clumps of gas \citep[e.g.,][]{vb06,vb07} and in others by interaction between gravitational instability in the outer disc and layered accretion in the inner disc \citep[e.g.,][]{zhu10}. Unfortunately it is not straightforward to draw direct comparisons between these results and ours, due to the complex effects of both the cooling and infall prescriptions used. We note, however, that the vertical shear effect observed in our simulations is intrinsically a three-dimensional phenomenon, and consequently cannot be observed in 2-D, vertically-integrated simulations. Our results point towards an additional mechanism for driving transient transport of angular momentum, and lend further weight to the well-established idea that low-order spiral modes drive accretion in protostellar discs. By contrast, to date only a handful of similar studies have been conducted in three dimensions. Most relevant here is the work of \citet{boley09} and \citet{kratter10}, who used three-dimensional hydrodynamics to study the formation and evolution of protostellar discs. \citet{boley09} used grid-based hydrodynamics with a similar set-up to that considered here: prescribed infall on to an already-present disc. In some cases the discs fragmented, while in others angular momentum transport was dominated by low-order spiral density waves. We note, however, that the discs in the simulations of \citet{boley09} are significantly more massive than ours ($q \sim 0.3$--0.5), increasing the importance of global modes. Given the additional differences between the simulations (most notably in the adopted cooling models) it is difficult to make direct comparisons, but in general our result -- that infall on to the disc enhances the importance of global modes -- seems consistent with those of \citet{boley09}. By contrast, in the models of \citet{kratter10} discs form and evolve in a self-consistent manner, and they were able to explore a larger range in parameter space than we have achieved here. They found, as we do, that low-order spiral modes dominate the transport of angular momentum, although this again may in part be driven by the fact that their discs are somewhat more massive than ours. However, \citet{kratter10} found that infall rates of $\gtrsim3$ times the disc accretion rate typically led to fragmentation, while we find that no fragmentation despite an infall rate nearly an order of magnitude higher than the ``local limit'' for disc accretion. Unfortunately it is not straightforward to compare these apparently contradictory results directly, due to the different prescriptions used for disc thermodynamics. \citet{kratter10} adopted an isothermal equation of state, and defined their models with two parameters (representing the accretion rate and angular momentum of the infalling gas). By contrast, we adopt an adiabatic equation of state with a parametrized cooling function, with the spherical envelope given a uniform initial temperature. The prescribed cooling time-scale is much longer than the dynamical time-scale (by a factor $\beta = 7.5$), so the infalling gas is effectively adiabatic. This results in slight heating of the infalling gas, and a corresponding increase in temperature in the disc. The increase in the disc temperature is not dramatic (see Fig.\ref{fig:panels_radprofile}), but given that the rate of spherical accretion scales as $c_s^3$ even this small difference could account for the factor of $\sim 3$ discrepancy between our results and those of \citet{kratter10}. Additional simulations, using different initial cloud temperatures and cooling laws, are required to investigate this issue in more detail, but such simulations are beyond the scope of this paper. It is not clear whether the isothermal or adiabatic approximation is more relevant to real discs; most probably both have some validity in different regions of the disc. We thus regard our results as complementary to those of \citet{kratter10}, and encourage further work in this area. \subsection{Applications to Observed Systems} Our results have obvious applications to the physics of star formation, in particular the formation of low-mass ($\sim 1$M$_{\odot}$) stars. Observations suggest that essentially all low-mass stars form with discs \citep[e.g.,][]{hll01}, and disc accretion is thought to play a major role in the build-up of stellar mass. Moreover, in the earliest, embedded phases gravitational instability is likely to be the dominant mechanism for angular momentum transport: such discs are insufficiently ionized to sustain transport via magnetohydrodyamic turbulence \citep{mt95,g96}, but both observations and theory suggest that they are indeed massive enough to be gravitationally unstable \citep[e.g.,][]{greaves08,andrews09,hueso05,vorobyov09}. Our results argue that accretion in such discs is likely to be highly transient, and in general terms are consistent with a picture where the bulk of the stellar mass is accreted during a small number of intense outbursts \citep[e.g.,][]{armitage01,lr05,vb06,vb10,zhu10}. The consequences of our results for massive star formation are less clear. Although discs are expected to form around massive, forming stars, observational evidence of their existence is somewhat thin \citep[e.g.,][]{cesaroni06,cesaroni07}. It is clear, however, that if such discs do indeed exist they are likely to be gravitationally unstable, but in this scenario the maximum stable disc accretion rate ($\lesssim 10^{-5}$M$_{\odot}$yr${^{-1}}$; \citealt{levin03,cesaroni06,rafikov07}) is much too low for these massive stars to accumulate their mass in a plausible time-scale. It has previously been suggested that the formation of massive stars is likely to be dominated by transient episodes and highly variable accretion \citep{cesaroni07}. Our results suggest that the global gravitational torques driven by infall on to the disc result in exactly this type of behaviour, and may be an important accretion mechanism in massive star formation. In addition, we suggest that our results may have important consequences for the formation of massive stars close to super-massive black holes. A large population of massive O- and Wolf-Rayet-type stars is now known to exist within $\sim 0.1$pc of the super-massive black hole (SMBH) at the centre of the Galaxy \citep{genzel03,ghez05}, and a popular scenario for the formation of these stars is ``in situ'' formation via the fragmentation of an accretion disc around the SMBH \citep{lb03,nayakshin06}. This picture has a number of attractive features, but an open question has always been how the stars attain their final masses. Both analytic theory and numerical simulations suggest that the initial fragment masses are small, $\sim 1$M$_{\odot}$, and that the bulk of the stellar mass is subsequently accreted from the SMBH disc \citep[e.g.,][]{rda08,br08}. However, the estimated infall rates through the Hill sphere on to these protostellar discs are extremely high, $\sim 10^{-4}$M$_{\odot}$yr$^{-1}$ \citep{ml04}, and in the local limit these discs are expected to fragment, preventing rapid growth of the protostars and limiting the resulting stellar masses \citep[see also][]{ml05,km06}. Our results suggest that discs subject to high infall rates may instead be able to transport angular momentum much more rapidly, through global modes of the GI, and this mechanism provides a potential solution to the ``accretion problem'' of massive star formation at the centre of the Galaxy.

1012.0724

1012

1012.5271_arXiv.txt

Neutrinos are allowed to mix and to oscillate among their flavor. Muon and tau in particular oscillate at largest values. Last Minos experiment claimed \cite{14} possible difference among their matter and anti-matter masses, leading to a first violation of the most believed CPT symmetry. Isotropically born atmospheric muon neutrino at $E_{\nu_{\mu}}\simeq 20-80$ GeV, while up-going, they might be partially suppressed by mixing in analogy to historical SuperKamiokande muon neutrino disappearance into tau, leading to large scale anisotropy signals. Here we show an independent muon rate foreseen in Deep Core based on observed SK signals extrapolated to DeepCore mass and its surrounding. Our rate prediction partially differ from previous ones. The $\nu_{\mu}$, $\bar{\nu_{\mu}}$ disappearance into $\nu_{\tau}$,$\bar{\nu_{\tau}}$ is leading to a ${\mu}$, $\bar{\mu}$ anisotropy in vertical up-going muon track: in particular along channel $3-5$ we expect a huge rate (tens of thousand of events) of neutral current events, charged current electron and inclined crossing muons. Moreover at channel $6-9$ we expect a severe suppression of the rate due to muon disappearance (in CPT conserved frame). Such an anisotropy might be partially tested by two-three string detection at $E_{\bar{\nu_{\mu}}}\geq 45$. A CPT violation may induce a more remarkable suppression of vertical up-going tracks because of larger $\bar{\nu_{\mu}}$ reduction for $E_{\bar{\nu_{\mu}}}\geq 35$.

The neutrino are very complex particles indeed. Their three light neutrino flavors mix in a complex way described by a matrix born only in last few decades \cite{12},\cite{18},\cite{19}. Their presence maybe recorded in small kiloton detector or larger (like SK and Super Kamiokande $22$ kiloton) ones.In largest size detectors as Icecube the higher characteristic $\nu_{\mu}$, $\bar{\nu_{\mu}}$ TeV energy do not oscillate much and they do not exhibit the negligible oscillation along our narrow Earth. However the new born Deep Core, while tuning to few tens or even a few GeV $\nu_{\mu}$, $\bar{\nu_{\mu}}$ energy, may hold memory of the $\nu_{\mu}$, $\bar{\nu_{\mu}}$ disappearance into tau. Many have foreseen the $\nu_{\mu}$, $\bar{\nu_{\mu}}$ disappearance in Deep Core \cite{09}\cite{15}. We did use their prediction to calibrate the eventual CPT violation influence into their future rate \cite{00}. Here we review these predictions and now we reconsider our preliminary estimate, based on Super Kamiokande ones \cite{03a}, estimates that partially disagree with the previous results \cite{09}\cite{15}\cite{10}as well as \cite{09a}. Deep Core, is a new telescope or better a counter event blurred at low energies (below $E_{\nu_{\mu}}\leq 30$ GeV) because muons are tracing tracks mostly projected along one string: the inner cone within $\sim 30^{o}$ may contain any neutrino arrival direction around the string axis azimuth angle . At higher energy $E_{\nu_{\mu}}\geq 45$ GeV the muon track, if inclined, may intersect two different strings leading to a much accurate (a few degree) angular resolution. Therefore the DeepCore may test different energy regions at different degree of resolution. Moreover most of the very inclined (respect the vertical) muons (below $E_{\nu_{\mu}}\leq 30$ GeV) may intersect briefly the one string leading to a few (four-five) detector optical module (DOM) signal, in a very short timing clustering. This mean that most of the inclined events may accumulate, as a noise, into the a few channel region that is at the same time the deposit of ten (or tens) GeV shower, mostly Neutral Current, (NC), event by $\nu_{\mu}$, $\bar{\nu_{\mu}}$,$\nu_{\tau}$, $\bar{\nu_{\tau}}$ and Charged Current, (CC), and NC due to electrons by $\nu_{e}$, $\bar{\nu_{e}}$. These few channel signals may also record rare (nearly a thousand) tau appearance. However the previous noise will make difficult to discover tau appearance, while muon disappearance is still viable. \subsection{SuperKamiokande rate versus DeepCore} The simplest way to estimate the Deep Core muon track and possible tau appearance (or better, the muon disappearance) has been shown by Icecube MC simulation \cite{10}, as described in figure \ref{123}, with our additional CPT violation expected influence\cite{00}. However we show here an independent derivation of the expected Deep Core rate, based on SuperKamiokande one \cite{01}. There are four main contribute to SK upgoing muons: Fully contained event (FC), born inside and decayed inside SK; partially contained event (PC), born inside but escaping outside the SK volume; the Upward Stopping Muons, born outside but decayed inside the detector; the upward through muon, just born outside , crossing and decaying again at external volume. Some care have been taken into account for the last Upward Stopping Muons and upward through muon: the SK location is deeply surrounded by mountain rock, while Deep Core is within much less dense ice. Therefore we suppressed the two last rates by the density ratio ($\simeq 2.6$) to calibrate the expected rate in DeepCore , amplified by extrapolated volume ratio ($\frac{V_{DeepCore}}{22kT}$) versus Deep Core one at each energy range. Indeed the Deep Core volume is variable with the muon energy values due to photodetector thresholds and muon Cherenkov luminosity. We considered here the preliminary Deep Core effective volume variability following last Icecube articles \cite{05},\cite{10},\cite{15},\cite{20}. Our result is described in the right side of figure \ref{123},\ref{05} in linear scale along the channel number. We assumed an averaged neutrino muon and anti-muon energy conversion, their length projected along the string at spread angle of $\theta\simeq 30^{o}$. The total event number derived by simplest SK-DeepCore translation is huge: $N \simeq 97.800$. Most of these events are not vertical but inclined. Therefore assuming a vertical beaming solid cone suppression (also to reconcile with Deep Core total expected rate) we selected only those events within a cone angle $\pm 33^{o}$, obtaining a fraction ($1- \cos\theta\simeq 0.16$) of the total rate, in this way now compatible with Deep Core preliminary global expectation $N \simeq 16.000$. \subsection{Rates and anisotropy} The calibrated muon rate figure \ref{05} shows in grouped channel graph number, the rate that we foresee following SK within a narrow vertical axis along each string. These prediction do not overlap with previous one. In particular as we did mention we foresee a huge rate of inclined events whose NC produce shower observable by $3-4-5$ channel: this very rough estimate is based on the NC, and electron CC,NC shower : they are well above $20.000$ NC (for $\nu_{\mu}$ $\nu_{\tau}$) with additional thousands of shower by CC and NC for $\nu_{e}$ and their antiparticles. The inclined tracks by nearly horizontal muons will excite the vertical string with a characteristic arrival time much shorter than any vertical shower event. Indeed the time difference in arrival for spherical shower along a string (each DOM at $7$ m separation) is nearly $\Delta t_{0}\simeq t_{0}= h/c = 23 ns$; by triangulation any horizontal muon tracks and its Cherenkov cone will record a quite shorter delay $\Delta t_{0}\simeq t_{0}\cdot(1-\frac{1}{\cos(\theta_{C}) \cdot n_{ice}}) \simeq 0.08 t_{0} = 1.84 ns$. Therefore this cluster of event, nearly coincident in time , might be a key test to calibrate the muon event rate in wide solid angle and possibly to meter the event rate at each channel. \begin{figure}[htb] \begin{center} \epsfig{file=dd.eps,scale=0.13} \epsfig{file=RateSK1.eps,scale=0.25} \epsfig{file=SK-Deep.eps,scale=0.22} \caption{Left: The expected event rate following \cite{05},\cite{10}, of atmospheric muon neutrinos whose tracks are detected along the photo-tubes (DOM) channel. The expected ($\nu_{\mu}\rightarrow\nu_{\tau} $) \emph{and the anti-neutrino} muon ($\overline{\nu}_{\mu}\rightarrow\overline{\nu}_{\tau} $) induce a suppression. Our prediction for CPT violation, following \cite{00} is also shown. The SK event rate in Deep Core extrapolated from SK for each different nature (FC,PC,Upward stopping, Through going) are shown in the center : dashed curve describe the upgoing stopping and upward through events; the FC and PC are described by lower continuous curve. The rate is much larger ($16$ times) probably because DeepCore is selecting only very vertical tracks, within $33^{o}$ from string axis, as shown in a normalized rate offered in figure \ref{05}.}\label{123} \end{center} \end{figure} \begin{figure}[htb] \begin{center} \epsfig{file=50GeV_new.eps,scale=0.25} \epsfig{file=Mixing.eps,scale=0.13} \caption{Left: In a CPT conserved (dashed lines) or in CPT violated case (continuous curve) scenario the neutrino conversion($\nu_{\mu}\rightarrow\nu_{\tau} $) \emph{and the anti-neutrino} muon ($\overline{\nu}_{\mu}\rightarrow\overline{\nu}_{\tau} $), as well as survival probability for given $50$ GeV energy at different distances \cite{00}. The thick vertical line shows the Earth diameter, the dashed ones the Cern-Opera,SK,icecube distances. The same oscillation in energy function and at earth diameter distance, in CPT conserved (dashed line) and CPT violation (continuous line) scenario reflects additional suppression of muon survival ($\overline{\nu}_{\mu}\rightarrow\overline{\nu}_{\mu} $) at energies $ E_{\nu} \geq 44$ GeV. }\label{34} \end{center} \end{figure} \begin{figure}[htb] \begin{center} \epsfig{file=SovraALL.eps,scale=0.24} \caption{The expected event rate at each muon (and anti-muon averaged) energy band, its corresponding muon (and anti-muon averaged) length, its corresponding channel number projected by a $\cos(\pi - \theta)$ factor; $\theta\simeq 33^{o}$. The rate are modulated by large red arrows standing for nearly one o two tens of thousand noise event, by a continuous black arrow due to CPT violated influence, by a broken line due to CPT conserved muon tau disappearance. The influence of dashed arrow (muons) is twice larger than anti-muon ones because the twice larger neutrino versus antineutrino cross section. See Fig. \ref{613}. }\label{05} \vspace{1cm} \epsfig{file=8GeV_2-3ch.eps,scale=0.22} \epsfig{file=13GeV_4ch.eps,scale=0.22} \epsfig{file=20GeV_6-6ch.eps,scale=0.22} \epsfig{file=32GeV_10-4ch.eps,scale=0.22} \epsfig{file=52GeV_17ch.eps,scale=0.18} \epsfig{file=82GeV_26ch.eps,scale=0.18} \epsfig{file=125GeV_40ch.eps,scale=0.18} \epsfig{file=200GeV_64ch.eps,scale=0.18} \caption{Eight different energy windows (corresponding to SK rate estimate above) for muon rates as a function of the arrival direction angle. Note the different suppression or enhancement of the oscillation in CPT violated (continuous) or conserved case (dashed) lines at vertical or quasi vertical arrival direction. These nearly vertical regions are above the vertical dashed line ($\pi - \theta \geq 2.6$ rad corresponding to $\theta \leq 33^{o}$) shown in each figure. The vertical downward North Pole points at $\theta \simeq \pi$ rad. The averaged energy label on the top of each figure corresponds to the channel region in figure above Fig. \ref{05} }\label{613} \end{center} \end{figure} \clearpage

We did show the rate of atmospheric muon neutrinos along Deep Core string channels with some comments on the expected anisotropy. The main results are: a) a very sharp peaked morphology in muon rate, namely a huge noise rate in low range $3-6$ channel (ten thousand or above event a year) followed by b), a deep minimum along channel $7-9$ due to muon disappearance (contrasted partially by eventual CPT violation), whose rate maybe below $2000$ event a year. c) A global decay and suppression of described events Fig (\ref{05}) in range above $10-50$, all along each channel, because of muon disappearance and additional anti-muon CPT violated suppression, by the vertical muon tracks that are more wide and they suffer larger disappearance leading to more anisotropic behavior respect to the averaged SK rates. The present rate estimate differ from the most known ones \cite{05},{10},{20}; however the CPT violation influence foreseen in our previous paper \cite{00} plays the same role: to reduce the common muon survival, making more anti-tau appearance, from channel above $\simeq 13$. The strong modulation by CPT violation at low channel ($3-6$) number,is quite remarkable, but it is nevertheless useless because a huge noise pollution by NC, electron CC,NC and nearly horizontal muon traces.

1012.5271

1012

1012.1586_arXiv.txt

We present results of a study of intermittency and multifractality of magnetic structures in solar active regions (ARs). Line-of-sight magnetograms for 214 ARs of different flare productivity observed at the center of the solar disk from January 1997 until December 2006 are utilized. Data from the Michelson Doppler Imager (MDI) instrument on-board the {\it Solar and Heliospheric Observatory} (SOHO) operating in the high resolution mode, the Big Bear Solar Observatory digital magnetograph and {\it Hinode} SOT/SP instrument were used. Intermittency spectra were derived via high-order structure functions and flatness functions. The flatness function exponent is a measure of the degree of intermittency. We found that the flatness function exponent at scales below approximately 10 Mm is correlated to the flare productivity (the correlation coefficient is - 0.63). {\it Hinode} data show that the intermittency regime is extended toward the small scales (below 2 Mm) as compared to the MDI data. The spectra of multifractality, derived from the structure functions and flatness functions, are found to be more broad for ARs of highest flare productivity as compared to that of low flare productivity. The magnetic structure of high-flaring ARs consists of a voluminous set of monofractals, and this set is much richer than that for low-flaring ARs. The results indicate relevance of the multifractal organization of the photospheric magnetic fields to the flaring activity. Strong intermittency observed in complex and high-flaring ARs is a hint that we observe a photospheric imprint of enhanced sub-photospheric dynamics.

1012.1586

1012

1012.1854_arXiv.txt

We report the results of observations of ten rotational transitions of water vapor toward the carbon-rich AGB (asymptotic giant branch) star IRC+10216 (CW Leonis), carried out with {\it Herschel}'s HIFI instrument. Each transition was securely detected by means of observations using the dual beam switch mode of HIFI. The measured line ratios imply that water vapor is present in the inner outflow at small distances ($\le \rm few \, \times \, 10^{14}\, cm$) from the star, confirming recent results reported by Decin et al.\ from observations with Herschel's PACS and SPIRE instruments. This finding definitively rules out the hypothesis that the observed water results from the vaporization of small icy objects in circular orbits. The origin of water within the dense C-rich envelope of IRC+10216 remains poorly understood. We derive upper limits on the H$_2^{17}$O/H$_2^{16}$O and H$_2^{18}$O/H$_2^{16}$O isotopic abundance ratios of $\sim 5 \times 10^{-3}$ (3 $\sigma$), providing additional constraints on models for the origin of the water vapor in IRC+10216.

The carbon-rich AGB (asymptotic giant branch) star IRC+10216 (CW Leonis) is surrounded by a dense, outflowing envelope that harbors a rich chemistry. Because of its proximity to Earth ($D \sim 100 - 170$~pc) and its large mass-loss rate ($\sim 1.5 - 3 \times 10^{-5} \rm M_\odot \, yr^{-1}$; see Crosas \& Menten 1997 and Sch{\"o}ier \& Olofsson 2000), IRC+10216 has proven to be the most valuable source available for the study of chemistry in a carbon-rich astrophysical environment; more than 60 molecules (e.g. He et al.\ 2008) have been detected within its circumstellar envelope (CSE), made up of the elements H, C, N, O, F, Na, Mg, Al, Si, S, P, F, and K. The photosphere of IRC+10216 has been significantly enriched in carbon -- owing to the dredge-up of helium-burning nuclear reaction products from the core -- leading to a large C/O ratio ($\sim 1.4$; Winters, Dominik, \& Sedlmayr 1994) and giving rise to a circumstellar chemistry that is qualitatively different from that typical of oxygen-rich AGB stars. In particular, while the most abundant molecules in O-rich CSEs are CO and H$_2$O, those in IRC+10216 are CO, HCN, and $\rm C_2H_2$. Prior to 2001, the only oxygen-rich molecules reported in the CSE of IRC+10216 were CO, SiO, and HCO$^+$ (Lucas \& Gu{\'e}lin 1999), but in the last decade four other O-bearing molecules have been detected: H$_2$O (Melnick et al.\ 2001), OH (Ford et al.\ 2003), $\rm H_2CO$ (Ford et al.\ 2004), and $\rm C_3O$ (Tenenbaum et al.\ 2006). The discovery of water vapor in IRC+10216, obtained by observations of the $1_{10}-1_{01}$ rotational transition near 557 GHz with the use of the {\it Submillimeter Wave Astronomy Satellite} (SWAS), was puzzling because models for thermochemical equilibrium in the photospheres of carbon-rich stars predict water abundances of only $\sim 10^{-10}$ relative to H$_2$ (e.g.\ Cherchneff 2006), many orders of magnitude below the detection threshold or the abundance $\sim 10^{-7}$ inferred (Ag{\'u}ndez and Cernicharo 2006; Gonz{\'a}lez-Alfonso, Neufeld and Melnick 2007) from SWAS observations. Several hypotheses have been proposed for the origin of the observed water vapor. Ford \& Neufeld (2001) proposed a model whereby the water vapor was released into the outflow by the vaporization of a Kuiper belt analog, in which orbiting icy objects -- heated by the enhanced luminosity of the AGB star -- underwent sublimation. Willacy (2004) subsequently suggested that the observed water could be formed from CO and H$_2$ by means of Fischer-Tropsch catalysis on metallic dust grains, whilst Ag{\'u}ndez and Cernicharo (2006; hereafter AC06) proposed that water might be formed in the outer envelope by means of radiative association of H$_2$ with O atoms produced by the photodissociation of $\rm ^{13}CO$ exposed to the ultraviolet interstellar radiation field (ISRF). Because SWAS had access to only a single transition of water vapor, it was unable to place observational constraints that could readily distinguish between these hypotheses. In a theoretical study (Gonz{\'a}lez-Alfonso, Neufeld and Melnick 2007; hereafter GNM) performed in anticipation of the {\it Herschel Space Observatory}, it was discussed how each of these hypotheses gave rise to a specific prediction for the spatial distribution of the emitting water vapor. GNM investigated how multitransition observations with {\it Herschel} might be used to determine that distribution, and showed, in particular, that the relative strength of lines of higher excitation than those accessible to SWAS was a decreasing function of the inner radius, $R_{\rm in}$, of the region containing water vapor. The vaporization of a Kuiper belt analog would lead to $R_{\rm in} \sim 2 \times 10^{15}{\rm \, cm} \sim 30 \,R_*$ (Model B in GNM; where $R_* \sim 8 \times 10^{13}\, \rm cm$ is the assumed stellar radius) because any icy object (at least of small size and on a circular orbit) within that radius would have been vaporized before IRC+10216 ascended the AGB. Water production by means of Fischer-Tropsch catalysis would result in a similar value of $R_{\rm in}$, according to calculations of Willacy (2004). The formation of water following the photodissociation of CO in an outer shell, as suggested by AC06, by contrast, would lead to a much larger $R_{\rm in} \sim 4.3 \times 10^{16} {\rm \, cm} \sim 500 \, R_*$ (Model C in GNM), and considerably smaller strengths for the higher-excitation water transitions. Conversely, were $R_{\rm in}$ significantly smaller than $2 \times 10^{15}\rm \, cm$ -- a possibility that had not been anticipated by any {\it specific} model -- then the high-excitation transitions would be relatively stronger (Model A in GNM). In this {\it Letter}, we report the results obtained from {\it Herschel} (Pilbratt et al.\ 2010) observations of 10 rotational transitions of water vapor, carried with the Heterodyne Instrument for the Far Infrared (HIFI; de Graauw et al.\ 2010). The observations and data reduction are described in \S 2, and the spectral line profiles and line intensities are presented in \S 3. In \S 4, we discuss the spatial distribution inferred for the water vapor in the CSE of IRC+10216, in the context of various hypotheses for its origin. We compare our results with those reported recently in an entirely independent study performed by Decin et al. (2010; hereafter D10) with the use of the PACS and SPIRE instruments on {\it Herschel}.

In Figure 3, we compare the measured line fluxes tabulated in Table 2 with the predictions of the GNM models. Here, the observed line fluxes are represented by black crosses for each of the ten detected transitions, ordered from left to right in increasing energy of the upper state, and the GNM predictions are shown by the solid lines for various values of the inner radius, $R_{\rm in}$, of the water emitting region. These predictions were ``calibrated" by adjusting the assumed water abundance to match the 557 GHz line flux detected by SWAS. Clearly, the best fit to the data is obtained for the smallest value of $R_{\rm in}$ considered by GNM, $4.5 \times 10^{14}\, \rm cm$ ($\sim 5.6 \, R_*$), but even for that value the model {\it underpredicts} the relative strengths of the transitions of highest excitation. Thus, our HIFI observations confirm the inference drawn by D10 -- from the detection of high-lying water rotational lines with PACS and SPIRE -- about the spatial distribution of water vapor; water is clearly present at distances smaller than $4.5 \times 10^{14}\, \rm cm$ ($\sim 5.6 \, R_*$) from the star. The presence of water that close to the star definitively falsifies the model proposed by Ford \& Neufeld (2001), in which the origin of the observed water vapor was the vaporization of icy comets in a Kuiper belt analog, because any icy object (at least of small size and on a circular orbit) within $4.5 \times 10^{14}\, \rm cm$ would have been vaporized before IRC+10216 ascended the AGB. The models proposed by Willacy (2004; i.e. Fischer-Tropsch catalysis) and particularly AC06 (production via radiative association in an outer shell) are similarly excluded. The fact that the 1113~GHz / 557~GHz line ratio is in good agreement with the GNM model implies that the ortho-to-para ratio is close to 3, the value assumed in the model. The dashed black line in Figure 3 shows our best-fit model, with $R_{\rm in}= 1.0 \times 10^{14}\, \rm cm$ ($\sim 1.3 \, R_*$), and a slightly larger water abundance $\rm H_2O/H_2 = 8.1 \times 10^{-8}$ than that assumed in GNM, now chosen to optimize the fit to {\it all} the transitions observed with HIFI. \begin{figure} \includegraphics[scale=0.50]{fig3.eps} \caption{Comparison of the measured water line fluxes (black crosses) with predictions of the GNM models for various values of the inner radius, $R_{\rm in}$, of the water-emitting region. The ten rotational transitions appear from left to right in order of increasing energy for the upper state (except for the 1153 GHz and 1097 GHz transitions which originate in the same upper state).} \end{figure} \begin{deluxetable}{ccccc}\tablewidth{0pt} \tablecaption{Water line fluxes measured toward IRC+10216} \tablehead{Transition & $\nu$ & $E_U$/k & $\int T_A dv$ & Flux$^{a,b}$ \\ & (GHz) & (K) & ($\rm K\,km\,s^{-1}$) & ($\rm 10^{-20}\,W \, cm^{-2}$)} \startdata $1_{10}-1_{01}$~(o)$^c$ & \phantom{1}556.936 & \phantom{1}26.7& 10.3 & 0.89\\ $2_{11}-2_{02}$~(p)\phantom{$^c$} & \phantom{1}752.033 & 136.9& \phantom{1}6.0 & 0.70\\ $2_{02}-1_{11}$~(p)\phantom{$^c$} & \phantom{1}987.927 & 100.8& \phantom{1}7.2 & 1.11\\ $3_{12}-3_{03}$~(o)\phantom{$^c$} & 1097.365 & 215.2 & 10.7 & 1.83\\ $1_{11}-0_{00}$~(p)\phantom{$^c$} & 1113.343 & \phantom{1}53.4 & 14.6 & 2.53 \\ $3_{12}-2_{21}$~(o)\phantom{$^c$} & 1153.127 & 215.2 & \phantom{1}3.4 & 0.62 \\ $3_{21}-3_{12}$~(o)\phantom{$^c$} & 1162.912 & 271.0 & \phantom{1}5.3 & 0.96 \\ $2_{21}-2_{12}$~(o)\phantom{$^c$} & 1661.008 & 159.9 & 15.5 & 4.09\\ $2_{12}-1_{01}$~(o)\phantom{$^c$} & 1669.905 & \phantom{1}80.1 & 35.1 & 9.56\\ $3_{03}-2_{12}$~(o)\phantom{$^c$} & 1716.770 & 162.5 & 24.2 & 6.74 \enddata \tablenotetext{a}{For an unresolved source at the beam center} \tablenotetext{b}{We conservatively estimate the flux uncertainty to be less than 15$\%$.} \tablenotetext{c}{The letters o and p indicate whether the transition is of ortho- or of para-water} \end{deluxetable} At present, the origin of water vapor in IRC+10216 remains poorly-understood. D10 have proposed an alternative model -- discussed in greater detail by Ag{\'u}ndez, Cernicharo and Gu{\'e}lin (2010; hereafter ACG) -- involving photochemistry in the inner envelope. As in the AC06 model, oxygen atoms are generated by photodissociation of $\rm ^{13}CO$ and SiO by the ultraviolet ISRF, but -- unlike the AC06 model -- the UV radiation is assumed to penetrate deeply into a clumpy CSE. In the D10/ACG model, the oxygen atoms are liberated close to the star, where the temperature is sufficient to drive H$_2$O production via a sequence of two H-atom extraction reactions with activation energy barriers: $\rm O(H_2,H)OH(H_2,H)H_2O$. Because the mean visual extinction through the CSE is $\sim 100$~mag, this scenario requires the existence of channels of greatly reduced extinction through which the ultraviolet radiation can penetrate. The consequences of this model for the abundances and spatial distribution of the many other molecules detected in IRC+10216, many of which have been observed interferometrically, has yet to be fully investigated. One possible test of this model might involve a search for H$_2^{17}$O or H$_2^{18}$O. Because the $\rm ^{12}CO$ photodissociation rate is sharply reduced by self-shielding, D10 have emphasized the importance of $\rm ^{13}CO$ as a source of atomic O that can react to form H$_2$O. The photodissociation rates for C$^{17}$O and C$^{18}$O would presumably be at least as large as that for $\rm ^{13}CO$. Thus, if $\rm ^{13}CO$, C$^{17}$O and C$^{18}$O were the only suppliers of atomic oxygen, the $\rm H_2^{16}O/H_2^{18}O$ and $\rm H_2^{16}O/H_2^{17}O$ ratios would approach the $\rm ^{13}C$O/C$^{18}$O and $\rm ^{13}C$O/C$^{17}$O ratios respectively. Given the $\rm ^{13}C/^{12}C$, $\rm ^{18}O/^{16}O$ and $\rm ^{17}O/^{16}O$ isotopic ratios determined by Kahane et al.\ (1992) and by Cernicharo, Gu{\'e}lin \& Kahane (2000) for IRC+10216, the $\rm ^{13}CO$/C$^{17}$O and $\rm ^{13}CO$/C$^{18}$O ratios are respectively $\sim$ 18 and 28, each a factor of $45$ (= $\rm ^{12}C/^{13}C$) smaller than the elemental $\rm ^{16}O$/$^{17}$O and $\rm ^{16}O$/$^{18}$O ratios in the CSE. If SiO or $\rm ^{12}CO$ are significant additional sources of atomic oxygen in the D10/ACG picture, then a $\rm H_2^{16}O$/H$_2^{17}$O ratio larger than 18 -- or a $\rm H_2^{16}O$/H$_2^{18}$O ratio larger than 28 -- might still be consistent with the model. Data obtained in a full HIFI spectral survey carried out toward IRC+10216 (Cernicharo et al.\ 2010) place upper limits on the flux of the H$_2^{17}$O (552.021 GHz) and H$_2^{18}$O (547.676 GHz) $1_{10}-1_{01}$ transitions. Comparing these with the observed flux in the 556.936~GHz H$_2^{16}$O $1_{10}-1_{01}$ transition, we determined that the spectral survey places 3~$\sigma$ lower limits on both the 556.936~GHz/552.021~GHz and 556.936~GHz/547.676~GHz line ratios of $\sim 100$. Taking account of optical depth effects, we find that these correspond to a lower limit of $\sim 200$ on both the $\rm H_2^{16}O/H_2^{17}O$ and $\rm H_2^{16}O/H_2^{18}O$ abundance ratios. Given the elemental isotopic ratios $\rm ^{16}O/^{18}O \sim 1260$ and $\rm ^{16}O/^{17}O \sim 840$ (Kahane et al.\ 1992), these limits imply that the abundance of the minor isotopologues could only be enhanced (by means of isotope-selective photodissociation of CO, for example) by at most a factor 4 (H$_2^{17}$O) or 6 (H$_2^{18}$O). Detailed modeling will be required to determine whether our non-detections of H$_2^{17}$O and H$_2^{18}$O are consistent with the water production mechanism proposed by D10/ACG.

1012.1854

1012

1012.4689_arXiv.txt

{Certain lines in spectra of the Galactic microquasar SS\,433, in particular the brilliant Balmer H$\alpha$ line, have been interpreted as emission from a circumbinary disk. In this interpretation the orbital speed of the glowing material is in excess of 200 km s$^{-1}$ and the mass of the binary system in excess of 40$M_\odot$. A very simple model for excitation of disk material is in remarkable agreement with the observations, yet it seems that the very existence of a circumbinary disk is regarded as controversial.} {To investigate whether analysis of optical data from H$\alpha$ and He I spectral lines in terms of a model, in which the disk is excited by radiation from the close environment of the compact object, can further illuminate the origin of these split spectral lines.} { A model in which the excitation of any given patch of putative circumbinary disk material is proportional to the inverse square of its instantaneous distance from the compact object was constructed and compared with published spectra, taken almost nightly over two orbital periods of the binary system. The H$\alpha$ and He I lines were analysed as superpositions of Gaussian components. } { The new model provides an excellent description of the observations. The variations of the H$\alpha$ and He I spectra with orbital phase are described quantitatively, provided the radius of the orbit of the emitting ring is not much greater than the radius of the closest stable circumbinary orbit. The observed variations with orbital phase are not consistent with an origin in a radially expanding ring.} { The new analysis has greatly strengthened the case for a circumbinary disk orbiting the SS 433 system with a speed of over 200 km s$^{-1}$ and presents supposed alternative explanations with major difficulties. If the circumbinary disk scenario is essentially correct, the mass of the binary system must exceed 40 $M_\odot$ and the compact object must be a rather massive stellar black hole. This possibility should be taken seriously.}

This paper is concerned with the origin of split lines in the optical spectrum of the Galactic microquasar SS 433, which have been attributed to a circumbinary disk. It contains a development of the model presented in Bowler (2010a), to which the reader is referred for the story so far and numerous relevant references. The H$\alpha$ and He I spectra here addressed are very usefully displayed in Schmidtobreick \& Blundell (2006), Fig.2. SS 433 is very luminous and unique in its continual ejection of plasma in two opposite jets at approximately one quarter the speed of light. The system is a 13 day binary and probably powered by supercritical accretion on the compact member from its companion. The orbital speed of the compact object is fairly well established but in order to determine its mass either a measurement of the orbital velocity of the companion is needed or a measure of the total mass of the system. Stationary emission lines in the spectra of SS 433 display a persistent two horn structure of just the kind expected for emission from an orbiting ring, or a disk, seen more or less edge on. The horn separation corresponds to a rotation speed in excess of 200 km s$^{-1}$ and attributed to material orbiting the centre of mass of the binary system implies a system mass in excess of 40 $M_\odot$. The H$\alpha$ spectra were originally discussed in Blundell, Bowler \& Schmidtobreick (2008). Departures from the pattern expected for a uniformly radiating ring are present and are more pronounced in He I emission lines. These departures were explained in terms of emission being stimulated by some kind of spotlight rotating with the binary, rather like the beam from a lighthouse (Bowler 2010a). That paper presented an analysis in terms of a very simple spotlight model; too simple to be realistic and yet yielding an astonishingly good description of the observations. Thus these observations fix rather well the mass of the SS 433 system and hence of the compact object, subject only to the proviso that the two horned structure in H$\alpha$ and He I is indeed produced in an orbiting circumbinary ring, perhaps the inner rim of a circumbinary disk. This notion has been greeted with some scepticism, one correspondent going so far as to describe it as "very speculative" and suggest a radially expanding ring as a more plausible explanation for the two horned structures (D. R. Gies, personal communication). My purpose here is to construct a realistic model of the appearance of spectral lines emitted from a disk excited by intense radiation from the environs of the compact object and compare these features both with the data and with such alternative mechanisms as have been suggested (although, so far as I am aware, no other suggested mechanism has been properly worked out). The model describes the data very well and its development has made it clear that the data do not agree with a radially expanding ring.

This analysis has greatly strengthened the case for a circumbinary disk generating emission spectra from material orbiting at approximately 250 km s$^{-1}$ at a radius close to $2A$. At the same time, it has placed formidable obstacles in the way of supposed alternatives. This is important because if the source is a circumbinary disk the accreting compact object must have a mass of somewhere in the region of 20 M$_{\odot}$ and so the compact object is probably a rather massive stellar black hole. The corollary is that the orbital speed of the companion in this binary must be in excess of about 130 km s$^{-1}$ (as reported by Cherepashchuk et al 2005), rather than the value extracted from absorption spectroscopy by Hillwig \& Gies (2008) and by Kubota et al (2010); 58 km s$^{-1}$. An extended campaign of absorption spectroscopy might be able to provide evidence for or against the suggestion that the origin of those absorption lines could be in orbiting circumbinary material (Bowler 2010a, 2010c). Perhaps my most important conclusion is that the case for the circumbinary ring is so strong that both observers and theorists should take seriously the possibility that the compact object is massive and not merely nudging neutron star territory. Fresh evidence, either for or against, is more likely to emerge if it is looked for.

1012.4689

1012

1012.3917_arXiv.txt

We present an analysis of the {\it Suzaku} observations of the prototype wind-blown bubble \NGC which is based both on use of standard spectral models and on a direct comparison of theoretical models with observations. The X-ray spectra of \NGC are soft and most of the X-rays are in the (0.3 - 1.5 keV) energy range. But, hard X-rays (1.5 - 4.0 keV) are also detected ($\sim 10$\% of the observed flux). The corresponding spectral fits require a relatively cool plasma with kT~$\leq 0.5$~keV but much hotter plasma with temperature kT~$\geq 2.0$~keV is needed to match the observed hard X-ray emission. We find no appreciable temperature variations within the hot bubble in \NGCE. The derived abundances (N, O, Ne) are consistent with those of the optical nebula. This indicates a common origin of the X-ray emitting gas and the outer cold shell: most of the X-ray plasma (having non-uniform spatial distribution: clumps) has flown into the hot bubble from the optical nebula. If the electron thermal conduction is efficient, this can naturally explain the relatively low plasma temperature of most of the X-ray emitting plasma. Alternatively, the hot bubble in \NGC will be adiabatic and the cold clumps are heated up to X-ray temperatures likely by energy exchange between the heavy particles (hot ions diffusing into the cold clumps).

Optical nebulosities trace large cavities around early-type stars (O, Of and Wolf-Rayet, WR) that are formed as the high-speed wind sweeps up ambient interstellar gas and compresses it into a thin shell. The flow pattern, resulting from this interaction, was first recognized by Pikelner (1968) and consists of two regions of shocked gas: one corresponding to the shocked stellar wind and the other to shocked interstellar gas. The two regions are separated by a contact discontinuity and the hot interior gas can cool via thermal conduction across the interface. Because of its high temperature, the interior of the wind-blown bubbles (WBB) is expected to emit in X-rays. In 1970’s and 1980’s, analytical solutions for the WBB structure were derived that revealed details in the physics of these interesting objects. For a review of the analytical works and the related physical ideas see Dyson (1981) and McCray (1983). It should be noted that the most complete (semi-)analytical study of the WBB structure was presented in Weaver et al. (1977). With the increasing computing power, the numerical hydrodynamic simulations became a powerful tool for studying the physics of these objects. The first numerical modeling of WBB was done by \citet{falle_75} and later on details of the adiabatic, radiative and conductive WBB were studied numerically (e.g., \citealt{ro_85}; Rozyczka \& Tenorio-Tagle, 1985a,b,c; Brighenti \& D'Ercole, 1995a,b, 1997; \citealt{de_92}; \citealt{gs_95}; \citealt{gs_96a}; \citealt{gs_96b}; \citealt{zhm_98}, 2000). Thanks to numerical simulations, it was found that the WBB are subject to numerous dynamic instabilities. These simulations also allowed a much more complete and complex physical picture be explored, namely, by following various phases of the WBB evolution when the wind of the central stars varies with time (for details see \citealt{gs_96a}; \citealt{gs_96b}). Observational properties of the optical nebula in WBB were described in detail in the early works of Chu et al. (1983 and the references therein) and Lozinskaya (1982; see also the book by Lozinskaya 1993 for an observational review). But, it should be emphasized that the physics of the hot interior (hot bubble) of a wind-blown bubble is a cornerstone in the entire physical picture of these objects. This is so since the hot bubble is the region where the stellar wind energy is stored and subsequently used for driving the entire structure. From such a point of view, X-ray observations of WBB are very important because they can provide us with details about the physical conditions in the hot bubble. The first successful X-ray detection of a WBB was that of \NGC by {\it Einstein} \citep{boch_88}. It was found that the X-ray emission from NGC 6888 is characterized by a plasma temperature kT $= 0.28 - 0.8$~keV (90\% confidence interval), or T $= 3.2 - 9$~MK. {\it ROSAT} observations were sensitive to cooler plasma and yielded a characteristic temperature T $\approx 2$~MK. They showed that the X-ray emission arises primarily in filament like structures \citep{wri_94}. {\it ASCA} observations suggested that even hotter plasma at T $\approx 8$~MK might be present in the bubble interior \citep{wri_05}. Preliminary results from a recent Chandra observation, {\it which covered only the northeast part} of \NGCE, detected plasma with a temperature of T $\approx 2-3$~MK and possible nitrogen enrichment \citep{chu_06}. Similar soft X-ray spectra were detected for another WBB, S308, with {\it ROSAT} \citep{wri_99} and {\it XMM-Newton} \citep{chu_03}. Unfortunately, none of these data provided us with solid grounds to draw a firm conclusion about the physical mechanism responsible for the X-ray emission from WBB. Obviously, the X-ray emitting plasma in the hot bubble has a moderate temperature (considerably lower than the one expected from a shock with velocity equal to that of the stellar wind) and the reason for this could be that the thermal conduction operates efficiently in the hot interior. But, it is necessary to have solid observational arguments that support or rule out such a theoretical expectation. This was the basic motivation for the recent {\it Suzaku} observations of \NGC which covered the entire object and provided us with the X-ray spectra having good photon statistics. We report here the results from our analysis of these data. The paper is organized as follows. We give some basic information about the WBB \NGC in Section \ref{sec:thebubble}. In Section \ref{sec:observations}, we briefly review the {\it Suzaku} observations. In Section \ref{sec:global}, we present results from the global spectral models. In Section \ref{sec:origin}, we discuss the origin of the hot gas in \NGCE. In Section \ref{sec:discussion}, we discuss our results and we present our conclusions in Section \ref{sec:conclusions}.

\label{sec:conclusions} In this work, we presented an analysis of the {\it Suzaku} data on \NGCE: the wind-blown bubble around the Wolf-Rayet star WR 136. For the first time, X-ray spectra of the entire nebula are obtained that have a very good quality allowing a detailed comparison between theory and observations. The basic results and conclusions are as follows. \begin{enumerate} \item The X-ray spectra of the WBB \NGC are soft and most of the emission is in the (0.3 - 1.5 keV) energy range although some emission ($\sim 10$\% of the observed flux) is found at higher energies (1.5 - 4.0 keV). The spectral fits require a relatively cool plasma with kT~$\leq 0.5$~keV but much hotter plasma with temperature kT~$\geq 2.0$~keV is needed to match the observed hard X-ray emission. We find no appreciable temperature variations within the hot bubble of \NGCE. \item One of the important results from the spectral fits is that the derived abundances (N, O, Ne) are very much the same as those of the optical nebula. Also, they are considerably different from the WN abundances \citep{vdh_86} assumed `typical' for the wind of the central star in the nebula. This indicates a common origin of the X-ray emitting gas and the outer cold shell, that is most of the X-ray plasma (likely concentrated in clumps as the X-ray images indicate) has flown into the hot bubble from the optical nebula. \item A direct comparison (in XSPEC) between 1D hydrodynamic models and the X-ray spectra of \NGC suggests a reduced mass-loss rate ($\sim 4-5\times10^{-6}$\dotM) of the central star in order to provide the correct value of the observed flux. We note that this figure is appreciably low for Wolf-Rayet stars and thus emphasizes the need of a global modeling of the entire system: central star, optical nebula and the hot bubble. \item The general physical picture that emerges from the current analysis is the following. Most of the X-rays come from clumps distributed in the hot bubble of \NGCE. These clumps originate from the cold optical nebula and they formed through various dynamic instabilities during the evolution of the wind-blown bubble. The shocked WN wind of the central star is likely the source of the very hot plasma (kT~$\geq 2$~keV) in this wind-blown bubble. If the electron thermal conduction is efficient, this can naturally explain the relatively low plasma temperature of most of the X-ray emitting plasma. Alternatively, the hot bubble in \NGC will be adiabatic and the cold clumps are heated up to X-ray temperatures likely by energy exchange between the heavy particles (e.g., a local diffusion of the rarefied hot plasma into the dense clumps). \end{enumerate}

1012.3917

1012

1012.1315_arXiv.txt

New \chan\ X-ray data and extensive optical spectroscopy, obtained with AAOmega on the 3.9 m Anglo-Australian Telescope, are used to study the complex merger taking place in the galaxy cluster Abell~2744. Combining our spectra with data from the literature provides a catalog of 1237 redshifts for extragalactic objects lying within 15\arcmin\ of the cluster center. From these, we confirm 343 cluster members projected within 3 Mpc of the cluster center. Combining positions and velocities, we identify two major substructures, corresponding to the remnants of two major subclusters. The new data are consistent with a post core passage, major merger taking place along an axis that is tilted well out of the plane of the sky, together with an interloping minor merger. Supporting this interpretation, the new X-ray data reveal enriched, low entropy gas from the core of the approaching, major subcluster, lying $\sim 2$\arcmin\ north of the cluster center, and a shock front to the southeast of the previously known bright, compact core associated with the receding subcluster. The X-ray morphology of the compact core is consistent with a Bullet-like cluster viewed from within $\sim 45^\circ$ of the merger axis. An X-ray peak $\sim 3$\arcmin\ northwest of the cluster center, with an associated cold front to the northeast and a trail of low entropy gas to the south, is interpreted as the remnant of an interloping minor merger taking place roughly in the plane of the sky. We infer approximate paths for the three merging components.

In a universe where structure grows hierarchically, clusters of galaxies are the latest structures to collapse and virialize \citep[e.g.,][]{springel2006}. There are three main modes of mass accretion onto rich clusters of galaxies: the steady infall of matter from the surrounding filamentary large scale structures, the discrete accretion of group-sized objects, and the extreme event of a major cluster-cluster merger. The latter are the most energetic events known in the Universe \citep{markevitch1998} and result in the violent reassembly of the cluster. Such a dramatic reconfiguration of the cluster results in a rapid change in the environment of its member galaxies, although what effect this has the galaxies themselves is yet to be fully understood. This is in part due to the complex nature of cluster mergers and, hence, the difficulty in obtaining a detailed picture of their properties. The well known Butcher-Oemler effect \citep[hereafter BO-effect;][]{butcher1978,butcher1984}, that the fraction of blue galaxies in the cores of rich clusters is significantly lower at the present epoch than it was $\sim$2.5\,Gyrs or more ago, reveals that a significant fraction of cluster galaxies have undergone rapid transformation in star-forming properties over the intervening period. Furthermore, the observed increase with redshift in the fraction of spiral galaxies in clusters at the expense of a commensurate decline in the S0 fraction \citep{dressler1997, fasano2000, desai2007,just2010} reveals a corresponding rapid evolution in galaxy morphology. While the BO effect appears to be widespread, the scatter in the blue galaxy fraction is large at all redshifts ($z>0.2$) and exceeds the uncertainties in the measurements \citep{butcher1984}. This indicates some internal cluster-specific mechanism is responsible for the scatter and \citet{margoniner2001} found that a large portion of this scatter can be attributed to the richness of a cluster, where less rich clusters have a higher blue fraction. Another intriguing possibility is that the scatter is driven by the hierarchical formation of clusters, in particular cluster mergers, which become more common with increasing redshift. Here it is quite conceivable that they would lead to an increased blue galaxy fraction through triggering star formation in the member galaxies \citep{kauffmann1995, metevier2000, miller2006}. In this context, spectroscopic observations of the Coma cluster have revealed a tantalizing correlation between the spatial distribution of post-starburst galaxies and regions involved in merger activity \citep{caldwell1993,caldwell1997,poggianti2004}. These observations strongly suggest that the merging process can affect the star-forming properties of the cluster galaxies. This spurred \citet{caldwell1997} to conduct a further study to search for ``abnormal'' spectrum galaxies, similar to those BO galaxies found at higher redshift, in four additional nearby clusters, three of which were selected on the basis of harboring clear substructure. Although their study was limited to galaxies of early-type morphology, the results indicated that there is a significant fraction of abnormal-type galaxies in nearby rich clusters and, most intriguingly, the triggering of starbursts leading to abnormal spectra occurs during the core passage phase of a merger. Recently, \citet{hwang2009} have provided further evidence for this scenario by comparing the galaxy properties in the post-core passage merger Abell~168 and the pre-core passage merger Abell~1750. Furthermore, radio observations have revealed that clusters which harbor evidence for major merger activity show an increase in radio activity amongst their galaxies \citep{miller2003,miller2006,venturi2001,venturi2000,johnstonhollitt2008}. However, there are counter-examples \citep{venturi2002}, suggesting that the merger phase is an important factor. According to the simulations of \citet{bekki2010}, this evolution may be driven by the significant increase in ICM pressure that a galaxy is exposed to during a cluster merger, in particular when the galaxy passes through regions affected by shocks during the core passage stage of the merger \citep[see also][]{roettiger1996}. {\it Therefore, an understanding of the merger dynamics and merging history of a cluster is critical when attempting to disentangle the effects leading to the triggering and halting of star formation in cluster galaxies.} The rich X-ray luminous cluster Abell~2744 at $z=0.3$ \citep[also known as AC118;][]{couch1984} is an interesting test case, being both a well known merging cluster and one that exhibits a significant BO effect. Two previous studies have addressed the dynamical state of the merger occurring in the core of Abell~2744. The first was the $Chandra$-based study of \citet[][hereafter KD04]{kempner2004}, who showed that Abell~2744 is undergoing a major merger in approximately the north-south direction. The second was conducted by \citet{boschin2006} who used 85 cluster member spectra to show that the merger has a significant velocity component along the line of sight, as suggested by \citetalias{kempner2004} \citep[see also; ][]{girardi2001}, and estimated the mass ratio of the merging subclusters to be $\sim 3:1$. This merger also quite naturally explains why Abell~2744 hosts one of the most luminous known radio haloes covering its central 1.8\,Mpc, as well as a large radio relic at a distance of about 2\,Mpc from the cluster center \citep{giovannini1999,govoni2001a,govoni2001b}. The presence of a significant blue galaxy excess---a blue fraction that was $2.2\pm0.3$ times that seen in the same core regions of nearby clusters---was first measured photometrically and confirmed spectroscopically by \citet{couch1987}, who found the cluster to have a blue galaxy fraction of $\sim 25$\%. Moreover, the spectroscopic study showed that the blue galaxy population in Abell~2744 was dominated by starburst and post-starburst galaxies. Abell~2744 therefore provides an excellent laboratory for studying the link between major merger activity and star formation activity in cluster galaxies. This paper reports the first step of this study, which is to determine more precisely the {\it phase} and {\it history} of Abell~2744's merger, in particular whether it is in a pre-core or post-core passage phase. As has been shown previously \citep[e.g.,][]{zabludoff1995,barrena2007,maurogordato2008,owers2009a,ma2009,ma2010}, the combination of high quality X-ray imaging and spectroscopy with comprehensive optical spectroscopy provides a powerful toolkit for disentangling the complex dynamics of cluster mergers. In this paper, we combine new \chan\ observations, which provide an additional 100ks of data to the existing 25ks analyzed in \citetalias{kempner2004}, with new AAT/AAOmega optical multi-object spectroscopy (MOS), which roughly doubles the number of spectroscopically confirmed cluster members over a 3\,Mpc radius region, and also doubles the number of cluster member spectra within the central regions studied by \citet{boschin2006}. These data sets are used both to detect substructure in Abell~2744, which is related to the merger activity, and to provide a more up-to-date interpretation of the merging history. The outline of this paper is as follows. In Section~\ref{chandra_data} we present the reduction and analysis of the new \chan\ data. In Section~\ref{optical_data} present details of the MOS observations and data reduction, the redshift catalog, and the cluster membership allocation procedure. In Section~\ref{substructure_detection} we present the methods used for substructure detection. In Section~\ref{SS_nat} we present our interpretations of the nature of the structures detected in the optical and X-ray data. In Section~\ref{merger_scen} we present our merger scenario based on the observations and interpretations presented in the paper. Finally, in Section~\ref{summary} we summarize our results. Throughout the paper, we assume a standard $\Lambda$CDM cosmology where $H_0 =70$\kms, $\Omega_m=0.3$ and $\Omega_{\Lambda}=0.7$. For the assumed cosmology and at the cluster redshift (z=0.3064), 1\arcsec=4.52\,kpc.

\label{summary} We have presented an analysis of the merging cluster Abell~2744 based on new \chan\ X-ray data and AAOmega optical spectroscopy. The deeper \chan\ data reveal a plethora of substructure including: \begin{enumerate} \item A significant structure to the north with a trail of X-ray emission to the south and curving westward. This northern core harbors cooler gas with higher metallicity than its surroundings, indicating it is a remnant cool core. \item The Bullet-like southern compact core which is surrounded by hot, shocked gas. There is an edge to the southeast of the southern compact core, which has the characteristics of a shock front with a Mach number $M=1.41_{-0.08}^{+0.13}$. \item A global peak in the X-ray emission with no corresponding galaxy overdensity, which is interpreted as the stripped atmospheres of the merging subclusters. \item The northwestern interloper has an edge to the northeast and a tail of emission extending towards the south. Analysis of the spectra of these features indicate that the edge is a cold front and that the tail of gas harbors cooler, lower entropy gas when compared to the surrounding gas. \end{enumerate} From the new AAOmega optical spectroscopy, we have found: \begin{enumerate} \item From 343 spectroscopically confirmed members within a 3\,Mpc radius, the redshift and velocity dispersion of Abell~2744 are, respectively, $z_{clus}=0.3064\pm 0.0004$ and $\sigma=1497\pm 47$\kms. \item The velocity distribution is significantly skewed and the KMM analysis reveals the velocity distribution in the central region is characterized by two distinct structures in velocity space, separated by $\sim 3000$\kms. \item Examination of the spatial distribution of the galaxies within the two partitions found in the KMM analysis of the velocity distribution reveal that for the negative peculiar velocity component the more significant galaxy concentration lies further to the north, approximately coincident with the cool, high metallicity northern core revealed by the \chan\ data. \item Applying the KMM algorithm to the combined velocity and spatial information, we find the statistically favored fit is one where the sample is partitioned into three. Two of the partitions correspond to the remnant cores of the two merging clusters in the central regions, while the third likely contains the galaxies which were in the outskirts of the two subclusters and are currently in the process of mixing to form the new cluster. The third substructure also contains those galaxies which belong to the northwestern interloper. \item The northwestern interloper detected in the \chan\ image is not detected as a dynamical substructure, nor as a substructure in local galaxy surface density when considering only spectroscopically confirmed members. An increase in local galaxy surface density is detected when considering cluster members defined using photometric redshifts, although only when including galaxies down to magnitude R=22.3. \end{enumerate} In combination, the observations presented here have been used to outline a skeleton hypothesis for the merger history of Abell~2744. In this scenario, we have identified two significant subclusters in the central region which have undergone a violent core passage and are now moving away from each other along a roughly north-south axis, with a large line-of-sight component. For the northwestern subcluster, we propose a scenario where the structure is traveling to the north/northeast after pericenter. Confirmation of this merger hypothesis requires detailed simulations and deeper optical spectroscopy in the central regions in order to probe the dynamics of the northwestern subcluster. The evidence for a post-core passage merger phase reported here, along with the significant enhancement of blue starburst and poststarburst galaxies residing within Abell~2744 \citep{couch1998}, provides further evidence that major cluster mergers can act as catalysts for rapid, environmentally driven evolution in the star forming properties of cluster galaxies and that this evolution occurs when the cluster galaxy environment is violently rearranged during the core passage phase of a major merger. Our large spectroscopic sample, in combination with deep radio data, can now be used to define a large sample of actively transforming galaxies which will further test this hypothesis of merger-induced galaxy transformation.

1012.1315

1012

1012.5700_arXiv.txt

We have made a new time-dependent calculation of the supernova production ratio of the long-lived isomeric state $^{180m}$Ta. Such a time-dependent solution is crucial for understanding the production and survival of this isotope. We include the explicit linking between the isomer and all known excited states. We have also calculated the properties of possible links to a conjectured excited state which might decrease the final isomer residual ratio. We find that the explicit time evolution of the synthesis of $^{180}$Ta using the available nuclear data avoids the overproduction relative to $^{138}$La for a $\nu$ process neutrino temperature of 4 MeV.

1012.5700

1012

1012.0549_arXiv.txt

{We discuss observational consequences of the curvaton scenario, which naturally appears in the context of the simplest model of chaotic inflation in supergravity. The non-gaussianity parameter $f_{\mathrm{NL}}$ in this scenario can take values in the observationally interesting range from $O(10)$ to $O(100)$. These values may be different in different parts of the universe. The regions where $f_{\mathrm{NL}}$ is particularly large form a curvaton web resembling a net of thick domain walls, strings, or global monopoles.}

One of the main reasons to introduce the curvaton scenario \cite% {Linde:1996gt,Enqvist:2001zp,Lyth:2001nq,Moroi:2001ct} was to obtain a realistic mechanism of generation of non-gaussian adiabatic perturbations of metric \cite{Linde:1996gt,Lyth:2002my}. Since that time, many interesting curvaton models were proposed. However, it would be nice to have a curvaton model which would be as simple as the basic chaotic inflation scenario with the potential $m^{2}\phi^{2}/2$ \cite{Linde:1983gd}. It would be good also to find a natural implementation of this scenario in the context of supergravity. This is the main goal of our work. In this paper we will describe a broad family of models of chaotic inflation in supergravity, which provide a natural realization of the curvaton theory. We will calculate the non-gaussianity parameter $f_{\mathrm{NL}}$ for the simplest versions of these models and show that this parameter takes different values in different parts of the universe, in agreement with the curvaton web scenario of Ref. \cite{Linde:2005yw}. If inflation is sufficiently long, the average value of $f_{\mathrm{NL}}$ in this scenario does not depend on the initial value of the curvaton field. We will also show that under certain conditions the parameter $f_{\mathrm{NL}}$ can take values in the observationally interesting range from $O(10)$ to $O(100)$.

In this paper we discussed the curvaton scenario, which naturally emerges in the simplest supergravity realization of the chaotic inflation scenario \cite% {Kawasaki:2000yn,Kallosh:2010ug,Kallosh:2010xz}. Investigation of this scenario consists of several parts. The main step is to find an average value of the curvaton field $\sigma $ after a long stage of inflation. One needs this to calculate the amplitude of perturbations of density of the curvaton field. We performed this investigation by analyzing the growth of the curvaton perturbations during inflation. To conclude this investigation, one should find the ratio $r$ of the energy of the curvaton field to the energy density of all other particles and fields at the time of the curvaton decay. This is a complicated and model-dependent problem, which requires study of reheating after inflation, the decay rate of the curvaton field, and the composition of matter at the time of the curvaton decay. In this paper, we simply treated $r$ as a free phenomenological parameter, but one should remember that all of the issues mentioned above should be addressed in a more detailed investigation. We analyzed the model with the simplest quadratic inflaton potential and with the curvaton mass given by $\alpha H^{2}+m^{2}$. Our investigation demonstrates that if inflation is long enough, then the average value of the curvaton contribution to the amplitude of metric perturbations, as well as the averaged value of the non-gaussianity parameter $f_{\mathrm{NL}}$, do not depend on initial conditions for the curvaton field. The final results depend on the inflaton mass $m$, and on the parameter $\alpha $, which is related to the curvature of the K{\"{a}}hler\thinspace\ manifold \cite% {Kallosh:2010xz}. However, the locally observable parameter $f_{\mathrm{NL}}$ and the amplitude of the curvaton perturbations may take different values in different parts of the universe and in certain cases they may significantly deviate from their averaged values \cite{Linde:2005yw}. Moreover, the average value of the parameter $f_{\rm NL}$ can be much greater than the value $f_{\rm NL}$ in the part of the universes with an average value of the field $\sigma$. For a certain choice of parameters, the value of the non-gaussianity parameter $f_{\mathrm{NL}}$ can be in the observationally interesting range from $O(10)$ to $O(100)$. The curvaton perturbations in our simple model have flat spectrum. This is a consequence of degeneracy of the masses of the inflaton and curvaton field at the end of inflation. One can change the spectral index by switching to a theory with a different inflaton potential. This can be easily realized in the new class of chaotic inflation models in supergravity, or by splitting the spectrum of fluctuations of the curvaton field into two branches with different masses \cite{Kallosh:2010ug,Kallosh:2010xz}. The last possibility can be realized by modifying the K{\"{a}}hler\thinspace\ potential, or by adding a term $\sim S^{3}$ to the superpotential. Another interesting possibility is to take the inflaton mass just a little bit smaller than $m \sim 6\times 10^{-6}$, to decrease the amplitude of the standard inflaton perturbations. Then one may compensate for this decrease by adding a small contribution of the curvaton fluctuations. This will result in a smaller amplitude of tensor modes and a larger spectral index $% n_{s}$, which would improve the agreement of the predictions of the simplest chaotic inflation models with the WMAP data. Also, as our calculations demonstrate, for certain values of parameters even a small contribution of the curvaton perturbations may dramatically increase the non-gaussianity of the combined spectrum of perturbations of metric. Thus, whereas the curvaton models are more complicated than the single-field inflationary models, they make the resulting scenario much more flexible, which may be important for a proper interpretation \cite{Easson:2010uw} of the coming observational data. Our final comment deals with the topological features of the distribution of perturbations in the curvaton scenario. We point out that in the theory of a single-component real curvaton field, the regions of the universe with large non-gaussianity form domain walls \cite{Linde:2005yw}, reminiscent of the exponentially thick cosmic domain walls. Meanwhile in the theory of a complex curvaton field, which was studied in the present paper, the regions of large non-gaussianity form exponentially thick cosmic strings. In more complicated theories, these regions may form separate islands of large local non-gaussianity, resembling global monopoles. Since these effects have a non-perturbative, topological origin, non-gaussianity in the curvaton scenario cannot be fully described by such tools as the familiar perturbation theory parameters $f_{\rm NL}$ and $g_{\rm NL}$.

1012.0549

1012

1012.5193_arXiv.txt

{The new VISual and Infrared Telescope for Astronomy (VISTA) has started operations. Over its first five years it will be collecting data for six Public Surveys, one of which is the near-infrared $YJK_\mathrm{s}$ VISTA survey of the Magellanic Clouds system (VMC). This survey comprises the Large Magellanic Cloud (LMC), the Small Magellanic Cloud, the Magellanic Bridge connecting the two galaxies and two fields in the Magellanic Stream.}{This paper provides an overview of the VMC survey strategy and presents first science results. The main goals of the VMC survey are the determination of the spatially-resolved star-formation history and the three-dimensional structure of the Magellanic system. The VMC survey is therefore designed to reach stars as faint as the oldest main sequence turn-off point and to constrain the mean magnitude of pulsating variables such as RR Lyrae stars and Cepheids. This paper focuses on observations of VMC fields in the LMC obtained between November $2009$ and March $2010$. These observations correspond to a completeness of $7$\% of the planned LMC fields.}{The VMC data are comprised of multi-epoch observations which are executed following specific time constraints. The data were reduced using the VISTA Data Flow System pipeline with source catalogues, including astrometric and photometric corrections, produced and made available via the VISTA Science Archive. The VMC data will be released to the astronomical community following the European Southern Observatory's Public Survey policy. The analysis of the data shows that the sensitivity in each wave band agrees with expectations. Uncertainties and completeness of the data are also derived.}{The first science results, aimed at assessing the scientific quality of the VMC data, include an overview of the distribution of stars in colour-magnitude and colour-colour diagrams, the detection of planetary nebulae and stellar clusters, and the $K_\mathrm{s}$ band light-curves of variable stars.}{The VMC survey represents a tremendous improvement, in spatial resolution and sensitivity, on previous panoramic observations of the Magellanic system in the near-infrared, providing a powerful complement to deep observations at other wavelengths.}

\label{intro} The cosmological paradigm for the formation and evolution of galaxies suggests that large structures formed as a sequence of mergers of smaller objects (White \& Frenk \cite{whi91}). The theoretical framework relies on cold dark matter simulations and is supported by high redshift observations (York et al. \cite{yor00}) and by studies of the cosmic microwave background (Spergel et al. \cite{spe03}), but the major difficulty is to reproduce the baryonic (stars, gas and dust) content of the Universe. Therefore, the study of the assembly process of nearby galaxies via resolved stars is a crucial aspect to understand how structures in the Universe form and evolve (Tolstoy et al. \cite{tol09}). In particular, dwarf irregular galaxies are well suited because their low metallicity and high gas content provide information about galaxies at an early stage of evolution. The closest prototypes of interacting dwarf galaxies that offer an excellent laboratory for this near-field cosmology are the Magellanic Clouds (MCs). The Magellanic system is located at a distance of $\sim57$ kpc (e.g. Cioni et al. \cite{cio00b}) and comprises: the Large Magellanic Cloud (LMC), the Small Magellanic Cloud (SMC), the Magellanic Bridge and the Magellanic Stream. The LMC is a dwarf irregular galaxy seen nearly face-on (e.g. van der Marel \& Cioni \cite{vdm01}) and sometimes referred to as a late-type spiral galaxy, rich in gas and actively forming stars. The SMC is a highly inclined dwarf irregular galaxy also referred to as a dwarf spheroidal galaxy (Zaritsky et al. \cite{zar00}) with less active star formation. The LMC is probably just a few kpc thick along the line-of-sight, but the SMC has a more complex structure that may extend up to 20 kpc along the line-of-sight (e.g. Westerlund \cite{wes97}, Groenewegen \cite{gro00}, Subramanian \& Subramaniam \cite{sub09n}). There is a bar embedded in each galaxy (Subramaniam \& Subramanian \cite{sub09m}, Gonidakis et al. \cite{god09}). The Magellanic system is metal-poor, the metallicity is about $1/2$, $1/4$ and $1/10$ that of the Sun for the LMC, SMC and the Bridge, respectively. The MCs have experienced an extended star formation history (e.g. Hill et al. \cite{hil00}; Zaritsky et al. \cite{zar02}, \cite{zar04}; Cole et al. \cite{col05}; Pomp\'{e}ia et al. \cite{pom08}; Gallart et al. \cite{gal08}; Carrera et al. \cite{car08}). The dynamical interaction between the MCs may be responsible for the various episodes of star formation (Zaritsky \& Harris \cite{zar04}) and for the creation of the Bridge (Irwin et al. \cite{irw85}, Gordon et al. \cite{gor09}) which connects the two galaxies and clearly has young stars associated with it (Irwin \cite{irw91}, Battinelli \& Demers \cite{bat98}). The Stream appears (to date) as a purely gaseous feature spanning more than $100$ deg in the Southern sky (Guhathakurta \& Reitzel \cite{guh98}). A tidal origin of the Stream from the interaction between the LMC and the Milky Way (MW) has been ruled out by new proper motion measurements (Kallivayalil et al. \cite{kal06a}, \cite{kal06b}). Alternative explanations are: ram pressure (Mastropietro et al. \cite{mas05}) and tidal origin from the interaction between the LMC and the SMC (Besla et al. \cite{bes10}). The interaction between the MCs and the MW is representative of the environmental effects that large galaxies with satellites (low-mass dwarf galaxies) experience elsewhere in the Universe. This suggests that the MCs may have entered the Local Group as part of an association (cf. Tully et al. \cite{tul06}, Moss \cite{mos06}, Knebe et al. \cite{kne06}) and that, in the future, a minor merger between the LMC and the MW may occur (cf. Ibata et al. \cite{iba94}, \cite{iba03}). An alternative is that the MCs may be tidal dwarfs expelled during a previous merger event involving M31 (Yang \& Hammer \cite{yan10}), although this would imply they have retained little dark matter from their parent halos (Barnes \& Hernquist \cite{bar92}). At present, basic assumptions are being challenged: What is the origin of the MCs? Do the MCs constitute a binary system and, if so, for how long? Have the MCs interacted with the MW or are they on their first approach (Besla et al. \cite{bes07}, D'Onghia \& Lake \cite{don08})? How have the star-formation histories of the LMC and the SMC been influenced by interaction? Does the geometry of the system depend on age and metallicity? How do star clusters form and evolve in the MCs? What is the fate of the MCs and will they merge with the MW? Will the Bridge evolve into a dwarf galaxy (Nishiyama et al. \cite{nis07})? Does the LMC have an ordinary bar and how does it influence the LMC evolution? Or, is the offset bar a separate galaxy being merged into the LMC disc? Does the LMC have a metal poor old halo? Why is there a significant difference in structure between the gas and stars in the SMC? Does the SMC have a bulge? To answer all these questions we must resolve the stellar populations and study them in detail. A fundamental step in this direction has been the many panoramic imaging surveys that have provided multi-wavelength observations of the Magellanic system. Except for the dedicated optical MCPS survey (Zaritsky et al. 2002, 2004) information about the overall population of the MCs has been obtained from surveys with different original goals, including microlensing optical surveys (e.g. MACHO - Alcock et al. \cite{alc00}, EROS - Tisserand et al. \cite{tis07}, and OGLE - Udalski et al. \cite{uda92}), and infrared sky surveys (e.g. IRAS - Schwering \cite{sch89}, 2MASS - Skrutskie et al. \cite{skr06}, and DENIS - Cioni et al. \cite{cio00a}). A continuation of OGLE is still in progress while other large-scale, near-infrared (near-IR) surveys (IRSF - Kato et al. \cite{kat07}) offer somewhat more sensitive data than 2MASS. The surveys in the mid-infrared with the {\em Spitzer Space Telescope} (SAGE - Meixner et al. \cite{mei06}, Bolatto et al. \cite{bol07}, Gordon et al. \cite{gor09}, Bonanos et al. \cite{bon10}) of the central part of the galaxies and an optical survey of the outermost regions (Saha et al. \cite{sah10}) have recently been completed. These surveys have provided data covering most of the electromagnetic spectrum, but their common depth is limited to moderately bright giant stars. The development of the VISual and Infrared Survey Telescope for Astronomy (VISTA - Emerson \& Sutherland \cite{eme10}) offers a unique opportunity to acquire near-IR data of unprecedented sensitivity in the Magellanic system. This is the underlying objective of the near-IR $Y J K_\mathrm{s}$ VISTA survey of the Magellanic Clouds system (VMC). This paper is organised as follows. Section \ref{prog} introduces the VMC survey and then describes the observing strategy and first observations. Section \ref{data} describes the data reduction steps for producing images and catalogues for individual observations, while Sect. \ref{vsa} describes the subsequent stages of reduction for deep and linked observations and presents the archival procedures. Section \ref{results} shows results from the first data. Section \ref{conclusions} concludes this study and the Appendix gives the coordinates of the VMC fields.

\label{conclusions} The VMC survey is a homogeneous and uniform $YJK_\mathrm{s}$ survey of $\sim$184 deg$^2$ across the Magellanic system (Fig. \ref{tiles}). It is an ESO Public Survey that started observations in November $2009$ and will run for approximately five years. The VMC survey parameters are described in Tab. \ref{vmcparam}. Images and catalogues will be delivered to the astronomical community at regular intervals with the first release expected in 2011. The VMC data will provide, among other things, a detailed history of star formation across the Magellanic system and a measurement of its 3D geometry. This paper presents the VMC survey strategy and first results aimed at assessing their scientific quality. These show the potential of the survey in addressing its main science goals and validates the expected sensitivity of the VMC data. Colour-magnitude and colour-colour diagrams show a wealth of substructures and a clear separation from Galactic foreground stars. These diagrams will form the core of the SFH analysis. To illustrate some of the scientific applications of the VMC survey, Cepheids and RR Lyrae stars are shown to display clear near-IR light-curves, PNe are detected thanks to the deep $K_\mathrm{s}$ band images dominated by Br$\gamma$ emission, and stellar cluster parameters are derived from best fit isochrones to the VMC colour-magnitude diagrams. The VMC survey will be of immense value to the astronomical community because the data will represent the only counterpart for existing optical surveys at a similar sensitivity (e.g. MCPS) and for the large number of unclassified objects observed by the {\em Spitzer Space Telescope} in the mid-IR (Blum et al. \cite{blu06}). Note that the near-IR 2MASS survey has observed only about 6\% of the stars that the VMC survey is expected to detect. The VMC data cover the bulk of the Magellanic system, as opposed to the tiny regions sampled by the Hubble Space Telescope, and the limited area covered by most of the other dedicated, ground-based observations at the same sensitivity. Among the other VISTA Public Surveys, the VISTA Hemisphere Survey (VHS) will also contribute to the investigation of the Magellanic system. The VHS is $\sim$3 mag shallower than the VMC survey but it covers an extended area around the Magellanic system to complement the VMC data.

1012.5193

1012

1012.3195_arXiv.txt

We present a general Bayesian formalism for the definition of Figures of Merit (FoMs) quantifying the scientific return of a future experiment. We introduce two new FoMs for future experiments based on their model selection capabilities, called {\em the decisiveness} of the experiment and the {\em expected strength of evidence}. We illustrate these by considering dark energy probes, and compare the relative merits of stage II, III and IV dark energy probes. We find that probes based on supernovae and on weak lensing perform rather better on model selection tasks than is indicated by their Fisher matrix FoM as defined by the Dark Energy Task Force. We argue that our ability to optimize future experiments for dark energy model selection goals is limited by our current uncertainty over the models and their parameters, which is ignored in the usual Fisher matrix forecasts. Our approach gives a more realistic assessment of the capabilities of future probes and can be applied in a variety of situations.

As cosmology becomes increasingly dominated by results emerging from large-scale observational programmes, it is imperative to be able to justify that resources are being deployed as effectively as possible. In recent years it has become standard to quantify the expected outcome of cosmological surveys to enable comparison, a procedure exemplified by the Figure of Merit (FoM) introduced by~\cite{Huterer_Turner} and later used by the dark energy task force (DETF) for dark energy surveys~\citep{Albrecht:2006um,Albrecht:2009ct}. Still in its infancy, however, is the topic of survey design, where an experiment is optimized, within design or cost constraints, to generate the best scientific outcome~\citep{Bassett05,BPN,Parkinson:2007cv,Parkinson:2009zi}. Both in quantifying and in optimizing survey capability, it is important to identify the scientific questions one hopes to answer. The DETF FoM measures the expected parameter constraints on a two-parameter dark energy model, using a Fisher matrix approach; this is an example of a parameter estimation FoM, in which the correct cosmological model is assumed to be known and the task is to estimate its parameter values (see also e.g.~\cite{Mortonson:2010px}). However, many of the most pressing questions in cosmology concern not parameters but models, i.e.\ the identification of the correct set of parameters to describe our Universe. Examples are whether cosmic acceleration is due to a cosmological constant, quintessence, or modified gravity, and whether or not the Universe has zero spatial curvature. These are model selection questions, hence forecasts of the capabilities of future probes should be assessed by their power to answer such questions, rather than the more limited question of the error they will be able to achieve assuming a given model is true (i.e., the usual Fisher Matrix forecast). Alternative FoMs, which quantify the ability of experiments to answer model selection problems, have been previously discussed by \cite{Mukherjee:2005tr}, \cite{Trotta:2007hy}, and \cite{BMIC_RTetal}.\footnote{For an alternative, essentially frequentist, perspective on this issue, see \cite{AmaraKitching}.} In this paper we present a comprehensive formalism for the construction of survey FoMs, incorporating both model and parameter uncertainty in light of the present observational situation. In order to do so, we build on the methodology introduced in~\cite{BMIC_RTetal}. We construct two new model selection FoMs, the {\em decisiveness} and the {\em expected strength of evidence}, which quantify the expected capability of an experiment to perform model comparison tests. For illustration we focus on the case of dark energy observations, though our formalism is broadly applicable.

We have presented a general Bayesian formalism for the definition of FoMs encapsulating the expected scientific return of a future experiments. Our method fully accounts for all source of uncertainties involved in the prediction, including present-day model and parameter uncertainties, and realization noise. It thus improves on the usual Fisher matrix methods by producing more realistic forecasts for the possible distribution of future experimental outcomes. We used this framework to define two Figures of Merit for probes that measure the dark energy equation of state in order to test the $\Lambda$CDM paradigm: the {\em decisiveness} $\Dec$ which quantifies the probability that a probe will deliver a decisive result in favour or against the cosmological constant, and the {\em expected strength of evidence} $\Expu$ that returns a measure of the expected power of a probe for model selection. We compared these quantities to the widely-used DETF FoM for a range of probes, and found that the rankings agree reasonably well, but that weak lensing and supernova probes have a higher than expected model selection power relative to their DETF FoM ranking. We also found, for our choice of prior, that there is a critical DETF FoM of around 70 below which probes are very unlikely to obtain a strong model selection result. An additional advantage of the formalism presented in this paper, and of any Figures of Merit that use it, is the possibility to include further observations, for example those that constrain the growth history or the presence of effective anisotropic stresses. One just extends the likelihood based on the predictions of the underlying models, but the procedure is unchanged, and the interpretation of the results is unchanged as well. There is therefore no need to define new FoM's as data analysis goals for future probes evolve. The methodology presented here is widely applicable to a variety of forecasting and optimization problems. Our application to the model selection capabilities of future dark energy missions is but a first step towards a fully Bayesian approach to performance forecast.

1012.3195

1012

1012.3640_arXiv.txt

We used four known chromospheric activity indicators to measure long-term activity variations in a sample of 23 M-dwarf stars from the HARPS planet search program. We compared the indices using weighted Pearson correlation coefficients and found that in general (i) the correlation between $S_{CaII}$ and \ion{Na}{i} is very strong and does not depend on the activity level of the stars, (ii) the correlation between our $S_{CaII}$ and H$\alpha$ seems to depend on the activity level of the stars, and (iii) there is no strong correlation between $S_{CaII}$ and \ion{He}{i} for these type of stars.

There is currently a focus on the search for planets orbiting around M-dwarf stars. Due to their low mass, it will be easier to find lower mass planets orbiting these stars with the radial-velocity (RV) technique. And therefore it is extremely important to access all sources of intrinsic noise that can degrade the quality of the detected RV signals. There are hints that the magnetic cycles of stars may induce RV capable of hiding the signals of extra-solar planets \citep[e.g.][]{dumusque2010} but not much is known about the long-term activity variations of M-dwarf stars. The Mt. Wilson survey showed that many solar-like stars have magnetic cycles similar to that of the Sun \citep{wilson1978,baliunas1995} but only one M dwarf was actually included in the sample. This star, HD\,95735, was shown to have long-term activity variations without a defined cycle. More recent studies however uncovered evidence for the existence of periodicity in the long-term activity of a few M stars. Our closest neighbour, Prox Centauri (dMe 5.5), was found to have a magnetic cycle with a $\sim$440 days period \citep{cincunegui2007a}. \citet{diaz2007b} also found a $\sim$760 days periodicity in the activity of the spectroscopic binary Gl375 (dMe 3.5). More recently, \citet{buccino2010} announced the detection of cycles with periods of $\sim$4 and $\sim$7 years for Gl299A (M1/2) and Gl752A (M2.5), respectively. In a recent paper, we studied the influence of the long-term activity cycles in the RV signals of a sample of 7 early-G and early-K stars known to have activity cycles \citep{santos2010}. We found no hints of RV induced variations by the activity cycles of these stars at the $\sim$1 m s$^{-1}$ level achieved by HARPS. In the present work we extend this study to the lower end of the main sequence by first analyzing the long-term behavior of the chromospheric activity indices and posteriorly compare them with the RV and parameters of the cross-correlation function of the stars.

We measured four activity indices for a sample of 23 M stars from the HARPS planet search program during a timespan of around 6 years. We compared the activity indices using weighted Pearson correlation coefficients and found that: \begin{itemize} \item There is a strong correlation between our $S_{CaII}$ and the \ion{Na}{i} indices. This confirms that the \ion{Na}{i} lines are good activity proxies for these cool stars as suggested by \citet{diaz2007a} for very active stars. \item As observed by \citet{cincunegui2007b} we found a great range of correlations between the $S_{CaII}$ and H$\alpha$ indices. Furthermore we found what appears to be a trend between the correlation and the average activity level of the stars as measured by the $S_{CaII}$ index. \item Although some authors suggest the use of the \ion{He}{i} line as a chromospheric activity proxy \citep[e.g.][]{saar1997b} we found that this index is not well correlated with $S_{CaII}$ for M dwarfs. \end{itemize} Since the signal-to-noise ratio in the \ion{Ca}{ii} H \& K lines is very low for M dwarfs, we suggest the use of the \ion{Na}{i} D1 and D2 lines, situated in a redder spectral region, as an alternative chromospheric indicator for this type of stars. These results may influence the way chromospheric activity is accessed in M-dwarf stars and contribute to the knowledge about the activity cycles of such stars. A more detailed study about this subject will be described in a future publication. Those results will then be used in the context of planet detection to search for trends between the long-term magnetic activity, RV, and parameters of the cross-correlation function in order to access at which level activity cycles might be inducing RV variations for these type of stars.

1012.3640

1012

1012.4917_arXiv.txt

The specific mechanism and astrophysical site for the production of half of the elements heavier than iron via rapid neutron capture (r-process) remains to be found. In order to reproduce the abundances of the solar system and of the old halo stars, at least two components are required: the heavy r-process nuclei ($A>130$) and the weak r-process which correspond to the lighter heavy nuclei ($A<130$). In this work, we present nucleosynthesis studies based on trajectories of hydrodynamical simulations for core-collapse supernovae and their subsequent neutrino-driven winds. We show that the weak r-process elements can be produced in neutrino-driven winds and we relate their abundances to the neutrino emission from the nascent neutron star. Based on the latest hydrodynamical simulations, heavy r-process elements cannot be synthesized in the neutrino-driven winds. However, by artificially increasing the wind entropy, elements up to $A=195$ can be made. In this way one can mimic the general behavior of an ejecta where the r-process occurs. We use this to study the impact of the nuclear physics input (nuclear masses, neutron capture cross sections, and beta-delayed neutron emission) and of the long-time dynamical evolution on the final abundances.

Half of the elements heavier than iron are produced by rapid neutron captures in a yet unknown astrophysical scenario. After the initial success of \cite{Woosley94} in reproducing observed solar r-process abundances, core-collapse supernovae and the subsequent neutrino-driven winds became one of the most promising candidates for the production of r-process elements because their extreme explosive conditions are very close to the ones needed for the r-process (see e.g.,~\cite{Hoffman97,Thompson.Burrows.Meyer:2001,Otsuki.Tagoshi.ea:2000}). Moreover, galactic chemical evolution models favor core-collapse supernovae, since they occur early and frequently enough to account for the abundances observed in old halo stars and in the solar system \cite{Ishimaru.Wanajo:1999,Ishimaru.etal:2004}. Although the necessary conditions to produce heavy elements ($A>130$) are identified \cite{Meyer92} (high entropies, low electron fractions, and short expansion timescales), these are not found in the most recent long-time supernova simulations \cite{Pruet.Hoffman.ea:2006,% arcones.janka.scheck:2007,Wanajo.Nomoto.ea:2009,% Fischer.etal:2010,Huedepohl.etal:2010}. When a supernova explodes, matter surrounding the proto-neutron star is heated by neutrinos and expands very fast reaching sometimes even supersonic velocity \cite{duncan.shapiro.wasserman:1986,Thompson.Burrows.Meyer:2001}. This neutrino-driven wind moves through the early supernova ejecta and eventually collides with it. The interaction of the wind with the slow-moving ejecta results in a wind termination shock or reverse shock where kinetic energy is transformed into internal energy. Therefore, the expansion velocity drops and the temperature (and thus the entropy) increases after the reverse shock. The matter near the proto-neutron star consists mainly of neutrons and protons due to the high temperatures in this region. When a mass element expands, its temperature decreases and neutrons and protons recombine to form alpha particles. The density also decreases but the triple-alpha reaction combined with different alpha capture reactions are still operating, resulting in heavy seed nuclei \cite{Woosley.Hoffman:1992,Witti.Janka.Takahashi:1994}. The evolution once the alpha particles start forming heavier nuclei depends on the neutron-to-seed ratio. The results presented here are based on our hydrodynamic simulations \cite{arcones.janka.scheck:2007} where the neutron-to-seed ratio is too low for the r-process to produce elements up to the third peak ($A=195$). In Sect.~\ref{sec:lepp} the nucleosynthesis obtained from such simulations is discussed. In Sect.~\ref{sec:rprocess} we have used the neutrino-driven wind simulations with the entropy artificially increased to study the impact of the long-time evolution and nuclear physics input on the dynamical r-process. More details can be found in \cite{Arcones.Montes:2010,Arcones.Martinez-Pinedo:2010}.

\label{sec:conclusions} Recent long-time supernova simulations do not produce r-process elements because the wind entropy is too low and the electron fraction high, even staying proton rich during several seconds \cite{Huedepohl.etal:2010}. However, the LEPP elements can be produced as we have shown by comparing for the fist time the LEPP pattern in UMP stars and integrated nucleosynthesis calculations based on hydrodynamical wind simulations \cite{Arcones.Montes:2010}. In proton-rich winds the LEPP pattern is very robust and reproduces observed abundances from UMP stars. Neutron-rich winds are necessary to explain the LEPP isotopes found in the solar system abundances, but they do not lead to a robust pattern and overproduced nuclei around A=90. This suggests that only a small fraction of the supernovae or of the mass ejected by them can be neutron rich. Future observations of isotopic abundances in ultra-metal poor stars could constrain the evolution of the electron fraction in the neutrino-driven winds and thus the neutrino properties (energy and luminosity). The impact of the long-time dynamical evolution and of nuclear masses on the r-process abundances can be still studied based on current simulations by artificially increasing the entropy. This mimics the hydrodynamical conditions of a neutrino-driven wind where the r-process does occur. We have found that the relevance of the different nuclear physics inputs depends on the long-time dynamical evolution \cite{Arcones.Martinez-Pinedo:2010}. If an $(n,\gamma)$-$(\gamma,n)$ equilibrium is reached, nuclear masses have a big influence on the final abundances. While for a cold r-process there is a competition between neutron capture and beta decay and these two process become relevant. This rises the importance of future experiments to measure nuclear masses that will provide a direct input for network calculations and constraints for the theoretical mass models. In both types of evolutions as matter decays to stability, our results show that neutron captures are key to understand the final abundances. Moreover, we found that beta-delayed neutron emission is important not only for the redistribution of matter, but also for the supply of neutrons. The late neutron captures are necessary to explain features in the solar system abundances, such as the rare earth peak. More experimental effort is necessary to test the validity of the current theoretical cross sections and more sensitivity studies of the impact of the neutron capture rates on the final abundances will give rise to new insights. \ack I would like to thank my collaborators, G.~Mart\'inez-Pinedo and F.~Montes, for the contributions to the work presented here. Support of Swiss National Science Foundation is acknowledged.

1012.4917

1012

1012.3706_arXiv.txt

We study the chaotic orbital evolution of planetary systems, focusing on {\it secular} (i.e., orbit-averaged) interactions, because these often dominate on long timescales. We first focus on the evolution of a test particle that is forced by multiple massive planets. To linear order in eccentricity and inclination, its orbit precesses with constant frequencies. But nonlinearities modify the frequencies, and can shift them into and out of secular resonance with the planets' eigenfrequencies, or with linear combinations of those frequencies. {\it The overlap of these nonlinear secular resonances drive secular chaos in planetary systems.} We quantify the resulting dynamics for the first time by calculating the locations and widths of nonlinear secular resonances. When results from both analytical calculations and numerical integrations are displayed together in a newly developed map, the ``map of the mean momenta'' (MMM), the agreement is excellent. This map is particularly revealing for non-coplanar planetary systems and demonstrates graphically that chaos emerges from overlapping secular resonances. We then apply this newfound understanding to Mercury. Previous numerical simulations have established that Mercury's orbit is chaotic, and that Mercury might even collide with Venus or the Sun. Guided by intuition from the test particle case, we show that Mercury's chaos is primarily caused by the overlap between resonances that are nonlinear combinations of four modes, the Jupiter-dominated eccentricity mode, the Venus-dominated inclination mode and Mercury's free eccentricity and inclination. Numerical integration of the Solar system indeed confirms that a slew of these resonant angles alternately librate and circulate. We are able to calculate the threshold for Mercury to become chaotic: Jupiter and Venus must have eccentricity and inclination of a few percent. Mercury appears to be perched on the threshold for chaos.

The question of the stability of planetary orbits in the Solar system has a long history, and has attracted the attention of some of the greatest scientists, including Newton, Laplace, Lagrange, Gauss, Poincar\'e, Kolmogorov, and Arnol'd. Newton thought that interplanetary perturbations are eventually destabilizing, and that divine intervention is required to restore the planets' orbits to their rightful places \citep{Laskar96}. Yet it is only over the last twenty years that the stability of the Solar system has been definitively settled, with the aid of computer simulations \citep{SussmanWisdom88, Laskar89, Quinn91,WisdomHolman,Lecar,Laskar09}. We now know that Newton was not far off: the Solar system is {\it marginally stable}: it is unstable, but on a timescale comparable to its age. In the inner Solar system, the planets' eccentricities chaotically diffuse on a billion-year timescale, with the two lightest planets, Mercury and Mars, experiencing particularly large variations. In fact, Mercury has roughly a $1\%$ chance of colliding with Venus or the Sun within the next five billion years \citep{Laskar09}. By comparison, the giant planets in the outer Solar system are well-spaced, and their orbital elements undergo largely quasiperiodic variations, exhibiting chaotic diffusion only on extremely long timescales \citep{Laskar96,murrayholman99}. The chaotic dynamics in the inner Solar system is primarily due to {\it secular} interactions \citep{Laskar08}. In general, interplanetary interactions can be decomposed into secular ones and MMR's (mean motion resonances---not to be confused with secular resonances). Secular interactions result from orbit-averaging the equations of motion. Since averaging a Keplerian orbit produces an elliptical ring, secular evolution can be thought of as interactions between elliptical rings. Secular timescales are long---they are longer than the orbital time by at least the ratio of the star's mass to that of a planet. By contrast, interactions driven by MMR's depend on orbital phase, and typically occur on the orbital timescale, or longer if some of the planets' orbital periods are close to integer ratios. Intuitively, one would expect that the dynamics on long timescales can be treated by averaging over the fast orbital phase---i.e., that they are secular in nature. This is true in the inner Solar system. It is also true more generally for well-spaced planets that do not happen to lie near mean motion resonances.\footnote{In the outer Solar system the dynamics is not mainly secular because the giant planets lie near a number of MMR's, such as the 5:2 between Jupiter and Saturn (the ``Great Inequality''), and the 2:1 between Uranus and Neptune. } {\it Linear} secular theory has been understood for hundreds of years, dating back to the famous solution of Laplace and Lagrange \citep[see][]{MD00}. To linear order in the planets' eccentricities and inclinations, secular theory reduces to a simple eigenvalue problem, with two eigenmodes per planet---one for the eccentricity degree of freedom, and one for the inclination. Each eigenmode has a constant amplitude and a longitude that precesses uniformly in time. But linear secular theory is clearly incapable of describing the chaotic orbits of the Solar system. Despite the importance of secular chaos, there has been surprisingly little theoretical understanding of it \citep[see, e.g., the review of Solar system chaos by][]{Lecar}. By contrast, chaos due to MMR's is well-understood, and accounts for the Kirkwood gaps in the asteroid belt \citep{wisdom83}, and for the very weak chaos of the outer Solar system planets, which is due to 3-body MMR's \citep{murrayholman99}. In all cases that have been studied in the Solar system, chaos is caused by overlapping resonances \citep{Chirikov,Lecar}. In linear secular theory, there can be secular resonances. And it is generally supposed that the chaos in the inner Solar system is caused by overlapping secular resonances. Yet thus far there has been little quantitative calculation. To our knowledge, the only previous theoretical work towards calculating secular chaos was by \cite{sid90}, who considered the coplanar case, as we describe below (\S \ref{sec:cop}). Numerical attempts to identify the mechanism of chaos in the inner Solar system were made by \cite{laskar90,laskar92} and \cite{SussmanWisdom}. These authors found that the angle associated with the (secular) frequency $(g_{\rm mercury}-g_{\rm jupiter})-(s_{\rm mercury}-s_{\rm venus})$ alternately librated and circulated in their simulations, where $g$ is the apsidal precession rate, and $s$ is the nodal precession rate (or, to be more precise, $g$ and $s$ here refer to the frequencies of the normal mode that is dominated by the corresponding planet). \cite{laskar92} also found that two angles associated with Earth and Mars, corresponding to $2(g_{\rm mars}-g_{\rm earth})-(s_{\rm mars}-s_{\rm earth})$ and $(g_{\rm mars}-g_{\rm earth})-(s_{\rm mars}-s_{\rm earth})$, alternately librated, and conjectured that the overlap of those secular resonances was responsible for chaos. But, as \cite{SussmanWisdom} note, Laskar's conjecture is not fully convincing, because there are too many unrelated angles that alternately circulate and librate, and it is not clear which are dynamically important. Furthermore, only one librating angle has been identified for Mercury. Yet chaos requires the overlap of at least two resonances, so why is Mercury chaotic \citep{Lecar}? Without a theory for secular chaos, the dynamics remain obscure. For example, why does instability in the Solar system occur at such low values of eccentricity and inclination ($\sim$ few percent)? What sets the timescale of the chaos? Can secular chaos shape the architecture of the inner Solar system \citep{Laskar96}? And can it shape the architecture of extrasolar planetary systems \citep{paper2}? Without a theory for secular chaos, we will be forever at the mercy of computer simulations. In this paper, we construct the theory for secular chaos of a test particle, and then apply the theory to Mercury. In \S \ref{sec:sem}, we present the test particle's equations of motion. In \S \ref{sec:cop}, we describe the coplanar solution, and in \S \ref{sec:iej} we generalize to the case when bodies have non-zero inclinations. In \S \ref{sec:realmercury}, we apply the theory to N-body simulations of the real Mercury. We conclude in \S \ref{sec:summary}.

\label{sec:summary} We have shown how secular chaos is driven by the overlap of secular resonances, both for a test particle modelled with a simplified Hamiltonian, and for the real Mercury. To linear order, secular frequencies are constant. But nonlinearities can shift planets into and out of secular resonance with each other, and when two resonances overlap, chaos results. In \S\S \ref{sec:sem}-\ref{sec:iej}, we focused on the evolution of a test particle in the presence of multiple massive planets. The test particle was evolved to leading nonlinear order, and the $e$'s, $i$'s, and precession rates of the planets (or more properly of the planet modes) were taken to be constant. We first considered the simple case with zero inclinations, as was first worked out by \cite{sid90}. In \S \ref{sec:iej}, we generalized to non-zero inclinations, when the test particle comes under the influence of one eccentric and one inclined planet mode. In that case, the particle has two free frequencies, its apsidal and nodal frequencies ($g$ and $s$). Each of these is altered by the particle's $e$ and $i$. Therefore each resonance traces out a one-dimensional curve in the particle's $e$-$i$ plane, or equivalently in its $p_e$-$p_i$ plane. A simple way to map out the dynamics is with the ``mean momentum map'' (MMM), whereby the particle's time-averaged $p_e$ and $p_i$ are plotted against each other for different initial conditions. This shows where the resonances are, how wide they are, and how their overlap leads to chaos (Figs. \ref{fig:low}-\ref{fig:pphi}). We calculated analytically the locations and widths of the four strongest resonances---the [1,0], [0,1], and [1,$\pm 1$]---and showed that these agreed with the numerical MMM results (Figs. \ref{fig:pplot}-\ref{fig:pplothi}). Chaos in this case emerges from the overlap of resonances of the form $[m,n]$ (eq. [\ref{eq:sig}]), with typically $n=m\pm 1$ and $m$ a small integer. This may be seen in the MMM, in surfaces of section (Fig. \ref{fig:sos}), and also by explicitly tracing the chain of librating angles (Fig. \ref{fig:resang}). We also examined the test particle's trajectories in Fourier space (Figs. \ref{fig:toy75}-\ref{fig:toyfftall}). In \S \ref{sec:realmercury}, we considered the orbital evolution of Mercury in N-body simulations. We showed that despite all the simplifications we made in the Hamiltonian models, the real Mercury behaved in a qualitatively similar manner. In particular: \bi \item Mercury's chaos is primarily driven by the $s_V$ and $g_J$ modes (i.e. the Venus- and Jupiter-dominated $i$ and $e$ modes), although the $g_V$ mode also plays a role. The nonlinear couplings between those modes and Mercury's own free modes (with frequencies $g_M$ and $s_M$) are primarily responsible for Mercury's chaos (Figs. \ref{fig:f75}-\ref{fig:fft1}). \item There are a slew of resonant angles that drive Mercury's chaos. Just as in the Hamiltonian model, a chain of resonant angles of the form $m(\pomega_M-\pomega_J)+n(\Omega_M-\Omega_V)$ show sequential librations, for integers $[m,n]$ (Fig. \ref{fig:rm}), where the angles refer to the phases of the free orbital elements. In addition, we identified a new four-angle combination, $(\pomega_M-\pomega_V)+(\Omega_M-\Omega_V)$, that can also undergo libration episodes (Fig. \ref{fig:tmerc}). \item Mercury is perched on the threshold of chaos. If one reduces the $e$'s and $i$'s of the planets by only 25\%, Mercury's motion becomes nearly regular (Fig. \ref{fig:f75}). This behavior is also apparent in the Hamiltonian model (Fig. \ref{fig:pphi}). We have also performed a $\kappa_{\rm nbody}=1.2$ simulation, in which the planets' initial $e$'s and $i$'s were increased $20\%$ (not shown). The result was violent instability, with Mercury ejected in $\sim 100$ Myr. \ei Having identified the secular resonances responsible for Mercury's chaos, and calculated the widths and locations of those resonances, we can calculate the threshold for Mercury's chaos. We do that here in an approximate way. First, since Mercury's apsidal and nodal frequencies differ from $g_J$ and $s_V$ by $\sim 20\%$, the co-ordinates in the $p_e$-$p_i$ plane where the two resonances overlap are around half that, or $p_{e**}\sim p_{i**}\sim 0.1$ (eq. [\ref{eq:pess}]). Second, the width of the chaotic overlap zone between the [1,-1] and the [1,0] (or [0,1]) is $\sim 2 (\epsilon_Ji_J)^{1/4}(p_{e**}p_{i**})^{1/8}$ (eq. [\ref{eq:over}]; we include here an extra factor of 2 to account for the difference between the half- and full-width, as described in \S \ref{sec:theory}) Therefore if $\epsilon_J\sim i_J\gtrsim p_{e**}^{3/2}/4\sim 0.01$, then the region of chaos will encroach upon the origin of the $p_e$-$p_i$ plane. This explains why Mercury can be chaotic even though the eccentricities and inclinations in the Solar system are at the level of a few percent. The work discussed in this paper can be extended in a number of directions. The theory can be extended to order-unity eccentricities and inclinations. Although that case will be more complicated, we suspect that the basic structure will remain, with resonant zones in the $e$-$i$ plane whose overlap leads to chaos. One can also attempt to build a theory that includes long-term diffusion and massive planets, as well as incorporating MMR's. A number of applications also come to mind, such as quantifying Mercury's chaotic diffusion and understanding how it came about that Mercury is perched on the threshold of chaos. The latter seems to be a clue for understanding how the Solar system arrived at its current marginally stable state. It would also be interesting to investigate the role of $\sigma'$ (eq. [\ref{eq:sigmapr}]). Is it an unnecessary coincidence for Mercury's chaos? Our theory can also be applied to Earth and Mars, for which librating angles have been identified that are similar to those we found for Mercury, i.e. angles of the form $m(\pomega_{\rm mars}-\pomega_{\rm earth})+ n(\Omega_{\rm mars}-\Omega_{\rm earth})$, with $[m,n]=[1,-1], [2,-1],$ and $[3,-2]$ \citep{laskar92,SussmanWisdom}. We also propose that secular chaos can play a role in shaping extra-solar planetary systems \citep{paper2}, and hence the theory of secular chaos might be applicable to extra-solar planets as well. \appendix

1012.3706

1012

1012.0533_arXiv.txt

It is evident from the results presented here that morphology remains a vital discipline in astronomical research. \small

1012.0533

1012

1012.2646_arXiv.txt

\setcounter{equation}{0} The late-time acceleration of the cosmic expansion invokes the mysterious dark energy or the modification of the gravity theory beyond Einstein's general relativity. Although there exist a number of dark energy and modified gravity models, both can be described by the effective equation of state (EOS) $\omega$ when one considers the expansion history of the Universe. Chevallier-Polarski-Linder (CPL) parametrization $\omega = \omega_0 + \omega_a (1 - a)$ is one of the suitable candidates to describe the general $\omega$ \cite{0009008,0208512}. However, models of different physical origins with the same background expansion history can not be separated with $\omega$. Thus, the growth of large scale structure is used as a complementary probe to segregate models \cite{0503644,0507184,0612452,07042421,07090307,08012431,08021068,08033292}. The growth index parameter (GIP) $\gamma$ defined by $\fr{d \ln \delta_{\m}(a)}{d \ln a} \equiv \Oma^{\gamma}$ is often used to discriminate different models using the linear matter perturbations \cite{Peebles1,Peebles2,Lahav}. Although $\gamma$ is time dependent by its definition even for the simple dark energy models \cite{09051522,09061643,09072108}, the constant value of it can be well matched with some dark energy and modified gravity models \cite{0507263,0608681,0701317,07101092,08024122,08081316}. However, it can be generalized as a time dependent $\gamma(a)$ \cite{9804015,07101510,08032236,09052470,09053444,09084379}. Behavior of $\gamma$ in the so-called Dvali-Gabadadze-Porrati (DGP) braneworld model \cite{0005016,0010186,0105068} has been widely investigated \cite{08033292,0701317,08024122,09051735,09053444,0401515}. $\gamma$ for a modification of gravitational action with a general function of the scalar curvature instead of the standard Einstein-Hilbert term, named f(R) gravity \cite{Buchdahl} has been also studied \cite{08032236,08093374,09035296,09082669}. In general, $\gamma$ has a scale dependence in this model. Also when one considers the scalar-tensor theories of gravity (STG) \cite{Bergmann,Nordtvedt,Wagoner}, one can obtain the specific $\gamma$ for the certain STG model \cite{07101510,0001066,08024196}. There is a useful approximate solution for the growth factor within general relativity \cite{0507263} \be \delta_{\m}(a) = a e^{\int_{0}^{a} (da'/a') [\Omega_{\m}(a')^{\gamma} - 1]} \, , \label{Linderg} \ee where $\Oma$ is the matter density contrast, $\delta_{\m} = \delta \rho_{\m} / \rho_{\m}$ is the linear density perturbation of the matter, and $\gamma = 0.55 + 0.05 [ 1 + \omega(z=1)]$ for $\omega > -1$. However, the accuracy of this solution is misinformed. As either $\omega_{0}$ or $\omega_{a}$ increases, the accuracy of the approximate solution given in Eq. (\ref{Linderg}) is decreased. For example, the error of approximate solution is about $1.2 \%$ when one consider ($\oo, \oa$) $=$ ($-0.78, 0.32$). If $\oa$ increases to $0.4$, then the error becomes $2.0 \%$ for the same value of $\oo$. There are both theoretical and phenomenological motivations for STG. The former is related to the existence a ubiquitous fundamental scalar coupled to gravity in theories which unify gravity with other interactions \cite{Fradkin,Callan,Lovelace,Green,Polchinski}. Also the dynamical equivalence between f(R) theories and a particular class of STG has been shown in the case of metric formalism \cite{Teyssandier,Magnano,9307034,0307338,09100434} as well as in the Palatini formalism \cite{0308111,0604028}. The later have several aspects. First, ``the lithium problem'' in the standard big bang nucleosynthesis (BBN) might be solved in STG due to the slower expansion than in general relativity before BBN, but faster during BBN \cite{0511693,0601299,08111845}. The weak lensing (WL) shear power spectrum in STG predicts the different one compared to in GR because they cause the different growth history of the matter \cite{0503644,0403654,0412120}. Integrated Sachs-Wolfe (ISW) effect probes modified gravities on cosmological scales through the matter potential relation \cite{9906066,08032238,SLAIP,09092045}. The crossing phantom $\omega < -1$ also can be naturally obtained in STG \cite{0504582,0606287,0610092}. There have been a number of reconstructions of specific STG models which is consistent with known observational constraints \cite{0011115,0107386,0508542,0612569,07053586,08031106,10031686,10061246,10112915}. However, we need to reconstruct theory without any specific theory {\it a priori}. Thus, we use both background and growth history parameters ($\omega, \gamma$) to find the viable subclasses of STG. We briefly review the basic background evolution equations of STG model in the next section. In Sec. 3, we also review the linear perturbation equations of the model. We derive the reconstruction equations for model functions $F(\phi)$ and $U(\phi)$ as a function of scale factor $a$ and check the viability of specific models in terms of parameters $\omega$ and $\gamma$ in Sec. 4. We conclude in Sec. 5. We also show the accuracy of the approximate solution in Eq. (\ref{Linderg}) for the general values of ($\oo, \oa$) and find the initial values of $\phi'$ and $U$ in the appendix.

\setcounter{equation}{0} Scalar-tensor gravities theories can produce many possible background evolutions which mimic dark energy models and other modified gravity models. We need to consider the growth of the linear matter perturbation to distinguish between models. However, STG models are strongly limited by the solar system constraint when we normalize the kinetic energy term $Z(\phi) = 1$. The main reason for this is that $\phi'$ becomes singular with this normalization. Thus, we may need to investigate the STG models with general $Z(\phi)$ in order to distinguish STG model with others. When we allow the time variation of the growth index parameter $\gamma$, usually the negative value of $\gamma_{a}$ models have the singular problem of $\phi'$. Some models with the positive $\gamma_{a}$ have the interesting features like large enough $F/F_0$ values at early epoch while mimic the dark energy models background evolution. However, these cases violate the solar system test and again this might be able to be cured when we consider STG models. If $\gamma_a$ is positive (negative), then $F(n)/F_0$ shows the convex (concave) shape with the minimum (maximum) as $1$. Thus, the value of $F(n)$ is bigger than $F_0$ in the past for the viable STG. There can be exception for this case, when $\gamma_a \simeq 0$ but positive. The main conclusion is that the viable STG models with $Z(\phi) = 1$ are not distinguishable from dark energy models or other modified gravity models when we strongly limit the solar system constraint. \appendix

1012.2646

1012

1012.5299_arXiv.txt

{ We describe a novel approach to accelerating Monte Carlo Markov Chains. Our focus is cosmological parameter estimation, but the algorithm is applicable to any problem for which the likelihood surface is a smooth function of the free parameters and computationally expensive to evaluate. We generate a high-order interpolating polynomial for the log-likelihood using the first points gathered by the Markov chains as a training set. This polynomial then accurately computes the majority of the likelihoods needed in the latter parts of the chains. We implement a simple version of this algorithm as a patch (\InterpMC) to \CosmoMC\ and show that it accelerates parameter estimatation by a factor of between two and four for well-converged chains. The current code is primarily intended as a ``proof of concept'', and we argue that there is considerable room for further performance gains. Unlike other approaches to accelerating parameter fits, we make no use of precomputed training sets or special choices of variables, and \InterpMC\ is almost entirely transparent to the user. }

\indent Cosmological parameter values are typically estimated using Monte Carlo Markov Chains [MCMC]. MCMC techniques are far more efficient than a brute force exploration of a parameter space with a realistic number of independent variables. Despite this, running Markov Chains for a broad range of parameter combinations -- standard procedure when analyzing the cosmological implications of a major astrophysical dataset -- remains computationally expensive. Consequently, algorithmic improvements that significantly increase the efficiency of this scheme without reducing its functionality are well worth exploring. Schematically, MCMC parameter estimation begins with a model, or {\em prior\/} which has a set of free parameters \cite{Christensen:2001gj,Knox:2001fz,RubinoMartin:2002rc,Lewis:2002ah,Verde:2003ey}. A likelihood function (derived for a specific combination of datasets) returns the relative probability that the ``observed sky'' was produced by the prior with a specific set of parameter values. After picking an initial point in the parameter space, we compute the likelihood at a new set of parameter values. The chain will {\em update\/} to this new point with probability, \begin{equation} \alpha(\Theta_n, \Theta_{n+1}) = \text{min} \left\{ 1, \frac{{\cal L}(\Theta_{n+1})q(\Theta_{n+1}, \Theta_n)}{{\cal L}(\Theta_n)q(\Theta_n, \Theta_{n+1})} \right\} \end{equation} where ${\cal L}(x)$ is the relative likelihood of parameter vector $x$ and $q(x, y)$ is the proposal density from $x$ to $y$. Within cosmology, \CosmoMC\ is a canonical and widely used implementation of the algorithm \cite{Lewis:2002ah}. A single likelihood can be computed in seconds, but a full set of chains requires the evaluation of many individual likelihoods. Computing ${\cal L}$ is a nontrivial task since we must generate the CMB angular power spectrum (or $C_\ell$) corresponding to our chosen parameter vector. Further, improvements in both angular resolution and signal-to-noise in future datasets will require $C_\ell$ to be computed with increased precision and over a larger range of $\ell$ than currently necessary, sharply increasing the computational cost. Hence there is a strong need to accelerate the MCMC analysis of cosmological data. A number of improvements to standard MCMC parameter estimation have been proposed to speed up or bypass the likelihood calculation. These include {\sc CMBFit} \cite{Sandvik:2003ii}, a polynomial fit to the WMAP 1-year likelihood, DASH \cite{Kaplinghat:2002mh}, which uses a combination of precomputation and analytic approximation, and WARP \cite{Jimenez:2004ct}, which combines interpolation with a careful choice of orthogonal variables \cite{Kaplinghat:2002mh} to accelerate the computation of the power spectrum. The most mature packages are {\sc CosmoNet} \cite{Auld:2007qz} which trains a neural network to provide likelihood values, and PICO \cite{Fendt:2006uh}, which uses a large, precomputed training set. MCMC estimation can be used with a vast range of problems. However, cosmological likelihoods -- particularly those derived from the Cosmic Microwave Background [CMB] experiments -- have two useful properties: they tend to be smooth functions\footnote{Exceptions to this rule certainly exist \cite{Easther:2004vq}, but it will not be a particularly burdensome restriction.} of the input parameters $x$ and evaluating them is computationally expensive. Consequently, we attempt to speed parameter estimates by caching computed values of the likelihood and then using interpolation to replace subsequent calls to the likelihood function within the MCMC code. The smoothness of the likelihood ensures that the interpolated likelihood will be a good approximation to the actual value. The computational overhead required by the caching and interpolation is small, compared to the cost of evaluating the likelihood directly, thus increasing the overall efficiency of the MCMC chains. Crucially, we make no use of precomputation when constructing our interpolation. Our technique essentially works because the MCMC analysis itself ``discards" enough information to reconstruct an interpolated likelihood; there is no need for a training set. This approach is straightforward computationally -- we use a stock interpolation routine and make relatively small changes to \CosmoMC\ itself. We implemented this algorithm as a patch to \CosmoMC, dubbed \InterpMC. We find that without any serious attempts to optimize the interpolation the runtime required for a typical parameter estimation is reduced by a factor of 2 to 4 for well-converged chains, with no degradation in the results. This improvement is nontrivial, although not as dramatic as that achieved by methods that rely on precomputation or analytic approximations. These can improve the runtime of MCMC code by an order of magnitude or more, but often at the cost of a good deal of extra work (whether analytic or computational), which must be performed beforehand. Further, \InterpMC\, works for any combination of datasets and cosmological model, is almost entirely transparent to the user, and offers a good deal of scope for future improvement. The source is available as a patch file to \CosmoMC.\footnote{{\tt http://easther.physics.yale.edu/interpmc.html} - at this point the code is offered as ``proof of concept", rather than production-ready code, but it has proved to be robust in a wide variety of settings.} The structure of this paper is as follows. In Section~\ref{sec:overview} we describe our algorithm, and the details of its implementation. Performance metrics and tests used to validate the interpolation scheme are discussed in Section~\ref{sec:test}. We test \InterpMC\ with a variety of datasets (WMAP, ground based CMB, supernovae, BAO) and scenarios, including both the usual concordance cosmology, and models with curvature, neutrinos, running and tensors, along with associated run-times. Finally, in Section~\ref{sec:disc} we summarize our work, and identify enhancements to \InterpMC\ that could further improve its performance.

} We describe an approach to reducing the runtime required for cosmological parameter estimation, based on interpolating cached likelihood values computed with a single set of MCMC chains. An implementation of this scheme, \InterpMC, is available as a patch to the standard MCMC package \CosmoMC. This analysis focuses on cosmological parameter estimation, but the fundamental approach is applicable to other problems where the likelihood is a smooth function of the free parameters, and computationally intensive to evaluate. We should ask {\em why\/} this approach works. Firstly, and unlike methods based on a precomputed training set (unless one has a very good idea of where to look, which is really only possible after an estimation has been performed), Markov Chains primarily evaluate points in regions of high likelihood, so our dataset is naturally weighted towards those points that will be frequently sampled during the estimation. Secondly, while MCMC estimates are much more efficient that a brute force search of the parameter space, they have no ``memory'' -- their behavior is determined solely by the current location of the chain in parameter space. Any interpolation scheme, whether generated on the fly or via a precomputed training set is, to some extent, exploiting the global properties of the likelihood function. In particular, cosmological likelihoods are typically smooth functions of the free parameters we are seeking to estimate. This greatly facilitates the construction of an interpolating polynomial, but this information is not exploited by the MCMC algorithm itself. As a side note, we would point that cosmological parameter estimation already makes substantial use of interpolation: CAMB directly computes only a subset of the underlying Fourier modes and the $C_\ell$ values and then fits the intervening points. Turning off this interpolation significantly increases the runtime of CAMB. Consequently, \InterpMC\ is extending the use of interpolation in cosmological parameter estimation, rather than introducing it for the first time. As presently implemented, the efficiency gain yielded by \InterpMC\ is typically between 2 and 4. Our primary goal here is simply to demonstrate the feasibility of this approach. We have deliberately employed conservative choices of the adjustable parameters in order to demonstrate that this approach works without careful tuning. Our experiments have shown that choosing more aggressive settings for the free parameters in \InterpMC\ can noticeably increases the efficiency gain delivered by \InterpMC\ without any modification to the underlying algorithm. Further, the interpolation requires a very small fraction of the total runtime, so even in the worst case scenario it cannot significantly slow \CosmoMC, relative to a standard run. We can see several specific algorithmic changes that may significantly improve the performance of \InterpMC. Firstly, we currently fit to all combinations of all free parameters, up to order $n$. However, many of these parameters are largely uncorrelated, and a more intelligent (but still automated) fitting procedure would focus on the subset of parameter combinations which make a nontrivial contribution to the fit, allowing the chains to begin making use of interpolated likelihoods at an earlier point in the run. Separately, the current interpolation is expressed in terms of the (normalized) free parameters in the chains. However, we can always choose a basis in which the free parameters are uncorrelated at second order, so that the coefficient of $x_i x_j$ in the interpolating polynomial is proportional to $\delta_{ij}$. In this case, the higher order coefficients (e.g. $c_{ijk} x_i x_j x_k$) would reflect the couplings between variables whose covariance vanishes at second order, and will likely have a smaller number of nontrivial terms than the corresponding polynomial written in terms of the unrotated variables. Further, we could move to a scenario in which some combinations of variables are interpolated at higher order than others. Finally, while we have focused on the ability to avoid precomputation, it would be simple to save the computed likelihoods to ``seed'' a future run\footnote{This might be useful in situations where the likelihood computation was unchanged, but the allowed parameter ranges were altered, or the chains were to be extended for better convergence.}, or to accelerate the subsequent computation of Bayesian evidence. Our goal here has been to interpolate the likelihood in a way that is essentially invisible to the user and does not modify the underlying MCMC algorithm in any way. However, we have effectively demonstrated that we can construct an accurate interpolation to the likelihood for a wide range of cosmological datasets and models, while expending less computational effort than that which is required for a full parameter estimation. The interpolating polynomial effectively yields the functional form of the likelihood in a substantial volume surrounding the peak. Consequently, the existence of a closed-form algebraic expression for the likelihood may allow us to pursue parameter estimation techniques that make direct use of this information, beyond standard MCMC techniques. It is clear that the computational cost of cosmological parameter estimation will continue to rise, particularly once data from Planck becomes available. Not only will this dataset include information at smaller angular scales than WMAP, the signal to noise will also be substantially improved, requiring a more accurate evaluation of the theoretical $C_\ell$, via CAMB. The computational cost of CAMB rises rapidly with the precision of the $C_\ell$, whereas the interpolation algorithm runs in negligible amounts of time. Further, both the likelihood code and CAMB itself undergo frequent modifications and updates, each of which would require a precomputed training set to be regenerated from scratch. In particular, the likelihood may undergo many changes as a new dataset is analyzed and reduced. Consequently \InterpMC's avoidance of trainings sets that are generated separately from the parameter estimation process guarantees that it will reduce the overall computational cost of parameter estimation, rather than just shifting it into the computation of the training set. To summarize, we have presented an initial implementation of an interpolation driven approach to accelerating MCMC cosmological parameter estimations, and shown that it can produce accurate results while reducing the resulting computational expenditure by at least a factor of two with a simple ``proof of concept'' implementation of this algorithm. Further, we have identified specific improvements to this approach that promise to substantially improve our current performance gains, and could be implemented in a production-ready version of \InterpMC.

1012.5299

1012

1012.2187_arXiv.txt

We study the cosmological inflation from the viewpoint of the moduli stabilization. We study the scenario that the superpotential has a large value during the inflation era enough to stabilize moduli, but it is small in the true vacuum. This scenario is discussed by using a simple model, one type of hybrid models.

Moduli fields play an important role in superstring theory. Couplings in 4D low-energy effective field theory are given as functions of vacuum expectation values (VEVs) of moduli fields. Thus, we need a stabilization mechanism of moduli VEVs. Indeed, moduli stabilization is one of important issues in string phenomenology and cosmology. The moduli potential has a small bump, which is related to the gravitino mass $m_{3/2}$, in many models of moduli stabilization, e.g. the Kachru-Kallosh-Linde-Trivedi (KKLT) potential \cite{Kachru:2003aw} and the racetrack potential \cite{racetrack}. When we consider the low-energy supersymmetry (SUSY) breaking scenario, the height of the above bump is quite low. That may lead to some problems. For example, we need the positive energy to derive inflational expansion of the Universe. If such a inflation energy is larger than this small bump of the moduli potential, the moduli would be destabilized and run away to infinity. Thus, we may have a constraint between the gravitino mass $m_{3/2}$ and the Hubble constant $H_{inf}$ during inflation, e.g. $m_{3/2} \geq H_{inf}$ in a simple model \cite{Kallosh:2004yh}. One way out is to use non-perturbative superpotential with positive exponents \cite{Abe:2005rx,Abe:2008xu,Badziak:2008gv}, although the KKLT superpotential and the usual racetrack superpotential include non-perturbative terms with negative exponents. However, the positive exponents are possible and have significant effects in the moduli stabilization. That is, the moduli potential with positive exponents could have a quite high barrier, which is independent of the gravitino mass. In this paper, we study another scenario to stabilize the moduli during the inflation era and to lead to low-energy SUSY. We consider the inflation scenario that the inflaton field $\phi$, which is different from moduli, drives the inflation dominantly. We assume that the would-be inflaton $\phi$ induces a large value of superpotential $\langle W(\phi) \rangle_{inf}$ during the inflation era. A similar idea has been studied in \cite{He:2010uk}. At any rate, such a large value of $\langle W(\phi) \rangle_{inf}$ could also induce a large mass of the modulus during the inflation era, and the modulus mass during the inflation would be of ${\cal O}(10)\times \langle W(\phi) \rangle_{inf}$, e.g. for the KKLT superpotential. Here we use the unit that $M_{Pl}=1$, where $M_{Pl}$ denotes the Planck scale. If such a mass is larger than the Hubble parameter $H_{inf}$ during the inflation era, the modulus is not be destabilized.\footnote{ This behavior is similar to the F-term uplifting scenario \cite{Saltman:2004sn,Dudas:2006gr,Abe:2006xp,Kallosh:2006dv}.} After the inflation ends, we assume that a small value of the gravitino mass $m_{3/2}$, i.e. $m_{3/2} \ll \langle W(\phi) \rangle_{inf}$, is realized at the potential minimum. At such a potential minimum, the modulus has a large mass compared with the gravitino mass $m_{3/2}$ such as ${\cal O}(10)\times m_{3/2}$ or more. Then, the modulus is stabilized at the true minimum. That is the realization of the F-term uplifting scenario at the potential minimum \cite{Saltman:2004sn,Dudas:2006gr,Abe:2006xp,Kallosh:2006dv}. To illustrate our scenario, we use the inflation model studied in Ref.~\cite{Nakai:2010km}, which is a kind of hybrid inflation models \cite{Dvali:1994ms,Dimopoulos:1997fv}. Its inflaton superpotential $W(\phi)$ is the Intrilligator-Seiberg-Shih (ISS) type of superpotential \cite{Intriligator:2006dd} with a deformation proposed in Ref.~\cite{Kitano:2006xg}. (See also for the inflation model with the ISS superpotential \cite{Craig:2008tv}.) With such a inflation superpotential $W(\phi)$, we study the modulus behavior during the inflation era for the KKLT superpotential and the racetrack superpotential including the case with a positive exponent. Indeed, similar superpotential models have been studied for inflation models \cite{hybrid-moduli}. However, our viewpoint differs from those. This paper is organized as follows. In section 2, we explain our scenario without explicit inflation superpotential. In section 3, we study our scenario by using an illustrating model. Section 4 is devoted to conclusion and discussion.

We have studied the inflation scenario from the viewpoint of the moduli stabilization. We have proposed the scenario that the superpotential has a large value during the inflation era and it induces a large mass of the modulus field. Then, we could realize the inflational expansion of the Universe with a sufficiently large value of e-folding $N_e$ without destabilizing the modulus. We study our scenario by using a simple model, the deformed ISS inflation model. It is found that in that model a sufficiently large e-folding can be obtained during the inflation, but at the final stage the modulus may be destabilized. Then, the modulus would run away to infinity in the simple KKLT superpotential. Thus, we would need another bump/barrier to avoid the runaway behavior of the modulus. Indeed, we have added the term with a positive exponent, although another type of potential bump may also be useful. As a result, this model shows a new aspect. The inflation is terminated when the modulus mass becomes tachyonic. That is, the modulus becomes the waterfall field. That has new and interesting effects on several cosmological aspects. We would study them elsewhere. At any rate, the modulus is not destabilized by the term with a positive exponent. Then, the system approaches toward the ISS vacuum. At the true (ISS) vacuum with the modulus stabilized, we can obtain the gravitino mass $m_{3/2}$, which is much smaller than the Hubble parameter $H_{inf}$ during the inflation. On the other hand, other inflation models would lead to a different behavior. The modulus might be stabilized in the whole inflation process in another inflation model with a large value of superpotential. It would be interesting to apply our studies to several types of inflation models. \subsection*{Acknowledgement} The authors would like to thank Y.~Nakai and O.~Seto for useful discussions. T.~K. is supported in part by the Grant-in-Aid for Scientific Research No.~20540266 and the Grant-in-Aid for the Global COE Program "The Next Generation of Physics, Spun from Universality and Emergence" from the Ministry of Education, Culture,Sports, Science and Technology of Japan.

1012.2187

1012

1012.2378_arXiv.txt

The advent of precise measurements of the cosmic microwave background (CMB) anisotropies has motivated correspondingly precise calculations of the cosmic recombination history. Cosmic recombination proceeds far out of equilibrium because of a ``bottleneck'' at the $n=2$ level of hydrogen: atoms can only reach the ground state via slow processes: two-photon decay or Lyman-$\alpha$ resonance escape. However, even a small primordial abundance of molecules could have a large effect on the interline opacity in the recombination epoch and lead to an additional route for hydrogen recombination. Therefore, this paper computes the abundance of the H$_2$ molecule during the cosmic recombination epoch. Hydrogen molecules in the ground electronic levels X$^1\Sigma^+_g$ can either form from the excited H$_2$ electronic levels B$^1\Sigma^+_u$ and C$^1\Pi_u$ or through the charged particles H$_2^+$, HeH$^+$ and H$^-$. We follow the transitions among all of these species, resolving the rotational and vibrational sub-levels. Since the energies of the X$^1\Sigma^+_g$--B$^1\Sigma^+_u$ (Lyman band) and X$^1\Sigma^+_g$--C$^1\Pi_u$ (Werner band) transitions are near the Lyman-$\alpha$ energy, the distortion of the CMB spectrum caused by escaped H~Lyman-line photons accelerates both the formation and the destruction of H$_2$ due to this channel relative to the thermal rates. This causes the populations of H$_2$ molecules in X$^1\Sigma^+_g$ energy levels to deviate from their thermal equilibrium abundances. We find that the resulting H$_2$ abundance is $10^{-17}$ at $z=1200$ and $10^{-13}$ at $z=800$, which is too small to have any significant influence on the recombination history.

The era of percent-level precision cosmology started with the exquisite measurements of the cosmic microwave background (CMB) anisotropies by the {\slshape Wilkinson Microwave Anisotropy Probe} ({\slshape WMAP}) satellite \cite{2003ApJS..148....1B}. The CMB anisotropies are a very useful tool for cosmologists for two reasons. First, the shapes and normalizations of the temperature and polarization spectra are sensitive to a host of cosmological parameters \cite{1996PhRvD..54.1332J}. Second, the physics underlying the CMB power spectrum is thought to be well understood. It can be calculated by linear perturbation theory of the Einstein and Boltzmann equations around a homogeneous, isotropic background \cite{1970ApJ...162..815P,1980PhRvD..22.1882B,1983ApJ...274..443B}; the perturbation equations can be solved rapidly by modern numerical codes that have achieved agreement at the 0.1\%\ level in code comparisons \cite{2003PhRvD..68h3507S}. However, to solve these equations one needs to know the number density of free electrons as a function of redshift $n_e(z)$, the so called recombination history, which enters into these equations through the Thompson scattering of photons from free electrons. The first cosmological recombination calculations were carried out more than 40 years ago \cite{1968ApJ...153....1P, 1968ZhETF..55..278Z}, showing the importance of non-equilibrium hydrogen recombination because of the high optical depth of the Lyman series lines in the early universe. A hydrogen atom can only reach its ground state from the H{\sc\,i} 2s level or 2p levels via two-photon decay and redshifting out of the Lyman-$\alpha$ line, respectively. The early analyses assumed Boltzmann equilibrium of all $n\ge 2$ levels of hydrogen, and thus had to follow only ionized hydrogen H$^++e^-$, excited hydrogen H$^\ast(n\ge2)$, and ground-state hydrogen H(1s). To obtain $n_e(z)$ to high accuracy, it is necessary to include additional physics. Thus theorists have considered helium recombination \cite{1969PThPh..42..219M, 1971PThPh..46..416M, 1995PhRvD..52.5498H,2005AstL...31..359D, 2007MNRAS.378L..39K, 2008PhRvD..77h3006S, 2008PhRvD..77h3007H, 2008PhRvD..77h3008S, 2008A&A...485..377R, 2010MNRAS.402.1221C}; deviations from Boltzmann equilibrium for the $n\ge 2$ levels of hydrogen \cite{2007MNRAS.374.1310C, 2010PhRvD..81h3005G, 2010MNRAS.407..599C, 2010MNRAS.407..658D, 2010PhRvD..82f3521A}; a host of two-photon processes \cite{2005AstL...31..359D, 2006A&A...446...39C, 2006AstL...32..795K, 2007MNRAS.375.1441W, 2008PhRvD..77h3007H, 2008A&A...480..629C, 2008PhRvD..78b3001H, 2010A&A...512A..53C}; the transport of photons near Lyman-$\alpha$ due to multiple resonant scattering \cite{1989ApJ...338..594K, 1990ApJ...353...21K, 1991Ap.....34..124G, 1994ApJ...427..603R, 2008AstL...34..439G, 2009A&A...496..619C, 2009PhRvD..80b3001H, 2009A&A...503..345C, 2010A&A...512A..53C}; and cross-talk among various lines and the photoionization continuum \cite{2007A&A...475..109C, 2008PhRvD..77h3006S, 2010PhRvD..81h3004K}. The workhorse recombination code {\sc Recfast}, used for {\slshape WMAP} parameter constraints, was a fitting function to such non-equilibrium calculations including all of the physical processes recognized as important in the year $\sim$2000 \cite{1999ApJ...523L...1S,2000ApJS..128..407S}, and there have been some subsequent updates \cite{2008MNRAS.386.1023W}. {\sc Recfast} was sufficiently accurate for the observations of its time, however to fully take advantage of the power of the {\slshape Planck} satellite data (launched 2009) it is important to find $n_e(z)$ to the sub-percent level \cite{2006MNRAS.373..561L,2008MNRAS.386.1023W}. This realization triggered a flurry of papers considering a host of new phenomena that could affect the recombination history to the percent and sub-percent level, culminating in two new publicly available codes that properly treat the radiative transfer effects in hydrogen and helium recombination \cite{2010arXiv1010.3631C, 2010arXiv1011.3758A}. The current paper is a continuation of the same effort to reach the required level of accuracy. We consider how the formation and destruction of hydrogen molecules (H$_2$) in the X$^1\Sigma_g^+$, B$^1\Sigma_u^+$ and C$^1\Pi_u$ electronic states can change the recombination history. The reason that H$_2$ might be able to change the recombination history is that the Lyman and Werner bands (X$^1\Sigma_g^+$--B$^1\Sigma_u^+$ and X$^1\Sigma_g^+$--C$^1\Pi_u$) are near the Lyman-$\alpha$ energy ($h\nu_{{\rm Ly}\alpha}=10.2\,$eV). Thus the excitation, de-excitation, photodissociation, and photoassociation of the H$_2$ molecule can shuffle photons between the red and the blue sides of the Lyman-$\alpha$ line. In an expanding Universe, a photon redder than Lyman-$\alpha$ is likely to simply redshift and eventually become a part of the far-infrared background, whereas a photon bluer than Lyman-$\alpha$ will redshift into the Lyman-$\alpha$ frequency and excite a ground-state hydrogen atom (which at $z>900$ would have been likely to be photoionized). At an order of magnitude level, one would expect this H$_2$-mediated redistribution to become possibly significant if the net optical depth in the Lyman and Werner bands $\tau_{\rm LW}$ were of order $10^{-3}$, which for the recombination-era density and Hubble rate, and total oscillator strengths of order unity, would require an abundance $x[{\rm H}_2]\equiv n({\rm H}_2)/n_{\rm H} \sim 10^{-12}$. Therefore we are interested in even tiny quantities of molecular hydrogen. It is worth noting that some models of early Universe chemistry have found $x[$H$_2]\gtrsim 10^{-12}$ \cite{1998A&A...335..403G, 2008A&A...490..521S} during the recombination epoch with simplified (i.e. not level-resolved) reaction networks. The calculation of the abundance of hydrogen molecules has already been considered by many authors {\em but for a different cosmological goal}, that is to assess the effect of the H$_2$ molecule on the cooling of metal-free gas and its implications for primordial star formation \cite{1967Natur.216..976S, 1968ApJ...154..891P,1969PThPh..41..835H, 1972PASJ...24...87Y, 1976ApJ...205..103H, 1984ApJ...280..465L, 1996ApJ...464..523H, 1996ApJ...467..522H, 1997ApJ...474....1T}. Since direct radiative association to the X$^1\Sigma_g^+$ electronic level, i.e. 2H$\rightarrow$H$_2($X$^1\Sigma_g^+)+\gamma$ is forbidden, two separate hydrogen atoms must reach the ground state of H$_2$ through an intermediate route. For the case of the post-recombination era when $k_{\rm B} T_{\text{CMB}} < 0.2\,$eV the accessible routes are through the H$_2^+$ \cite{1967Natur.216..976S} and H$^-$ \cite{1968ApJ...154..891P,1969PThPh..41..835H} intermediate states. Indeed, complex reaction networks have been constructed to follow hydrogen chemistry \cite{1984ApJ...280..465L, 1998A&A...335..403G, 1998ApJ...509....1S, 2000MNRAS.316..901F, 2002JPhB...35R..57L, 2006MNRAS.372.1175H}. These have identified in particular the significance of the recombination-induced CMB spectral distortion \cite{2006A&A...458L..29C, 2008A&A...485..377R, 2008A&A...488..861C} in controlling pregalactic photochemistry, the importance of rate coefficients \cite{2006ApJ...640..553G, 2008MNRAS.388.1627G, 2009MNRAS.393..911G}, and the importance of following transitions among the various rotational and vibrational levels of the H$_2^+$ ion at $z<500$ \cite{2006MNRAS.372.1175H}. However, at the redshift of interest for this paper, $z \sim 1000$, there is another route for the formation of hydrogen molecules: the inverse Solomon process \cite{2006ApJ...646L..91D}. At this era there are enough ultraviolet photons (both blackbody photons and spectral distortion photons) to facilitate the photo-attachment of two hydrogen atoms into an excited H$_2$ molecule in one of the rovibrational levels of either the B$^1\Sigma^+_u$ (Lyman band) or C$^1 \Pi_u$ (Werner band) electronic states with energies $\sim 10\,$eV. The excited H$_2$ molecule will re-emit the photon and decay to either a bound H$_2$(X$^1\Sigma_g^+$) molecule, or to the continuum of the X level (i.e. to two H atoms). In equation form, \begin{equation} 2{\rm H} + \gamma \leftrightarrow {\rm H}_2({\rm B}^1\Sigma_u^+,{\rm C}^1\Pi_u) \leftrightarrow {\rm H}_2({\rm X}^1\Sigma_g^+) + \gamma. \label{eq:H2process} \end{equation} This mechanism and the charged-particle processes (H$^-$, H$_2^+$, and HeH$^+$) control the H$_2$ abundance at high redshift. The possible effect on hydrogen atom recombination is the main focus of this paper. We note that Ref.~\cite{2006ApJ...646L..91D} found only a small production of H$_2$ via this mechanism, but they did not include the spectral distortion photons in their rate coefficient and hence the {\em total} rate of H$_2$ production could be many orders of magnitude larger. Of course, the same spectral distortion also drives H$_2$ photodissociation -- the left arrows in Eq.~(\ref{eq:H2process}) -- so the net effect on the H$_2$ abundance requires a detailed calculation. Deviations from thermal equilibrium abundances arise not from the amplitude of the ultraviolet photon spectrum, but the way in which its peculiar shape beats against the forest of H$_2$ lines and dissociation continua. Since we work at $z>800$ we will not distinguish the matter versus radiation temperature in this paper. This paper is organized as follows: in Sec.~\ref{sec:AC} we write down the rate equations for the bound-bound and bound-free transitions. These equations are then solved in the steady state approximation and the results are presented in Sec.~\ref{sec:R}. We discuss the size of the H$_2$ abundances found in the previous chapter on the absorption of Ly-$\alpha$ photons and conclude in Sec.~\ref{sec:DandC}.

\label{sec:DandC} Since the Lyman and Werner band transition energies of H$_2$ are near the H Ly$\alpha$ energy, it is expected that the abundances of H$_2$ energy levels deviate appreciably from their thermal abundances. This is because the photon phase space density has been distorted by the redshifted Ly$\alpha$ photons as in the standard hydrogen atom recombination picture. However, it is not clear from the outset whether this distortion to the photon phase space density increases or decreases the abundances of H$_2$ levels compared to their thermal abundances, since the spectral distortion photons accelerate both the production and destruction of H$_2$. To answer this question and ultimately to see to what extent the H$_2$ molecules can affect the recombination history we have in this paper carried out a detailed calculation including all of the rovibrational levels of the H$_2$ X$^1\Sigma_g^+$, B$^1\Sigma_u^+$, and C$^1\Pi_u$ electronic states up to rotational number $J=20$, together with the charged species relevant to the formation of hydrogen molecules, that is H$_2^+$, HeH$^+$ and H$^-$. We have calculated the bound-bound and bound-free dipole transition rates for the Lyman and Werner bands of the hydrogen molecule using the Born-Oppenheimer approximation. Special care has been taken to find the resonances of the bound-free transitions. The rate equations connecting the energy levels are then solved in the steady state approximation and the level abundances are found by a matrix inversion for each given redshift. The main result of our paper is that the shape of the CMB spectral distortion reduces the abundance of H$_2$ compared to the thermal abundance, resulting in low H$_2$ abundances throughout the recombination epoch; see Fig.~\ref{fig:sum_x}. The inclusion of the quadrupole transitions among rovibrational levels of the X electronic state increases the H$_2$ abundance, and adding the charged particle processes increases the H$_2$ abundances yet more, while remaining below the thermal abundance. We find $x[$H$_2]\sim 10^{-16}$ during most of the recombination epoch, rising to $10^{-13}$ at $z=800$. We conclude that -- despite the high cross section for Lyman and Werner band absorption -- H$_2$ is not relevant for determination of the primordial recombination history and CMB anisotropies.

1012.2378

1012

1012.2839_arXiv.txt

Supernovae (SNe) driven winds are widely thought to be very influential in the high-redshift Universe, shaping the properties of the circum-galactic medium, enriching the intergalactic medium (IGM) with metals and driving the evolution of low-mass galaxies. However, it is not yet fully understood how SNe driven winds interact with their surroundings in a cosmological context, nor is it clear whether they are able to significantly impact the evolution of low-mass galaxies from which they originate by altering the amount of cold material these accrete from the cosmic web. Indeed, due to the strong constraints on resolution imposed by limited computational power, all cosmological hydrodynamics simulations to date resort to implementing more or less physically well motivated and complex subgrid models to trigger galactic winds. To explore this issue, we implement a standard Taylor-Sedov type solution, widely used in the community to depict the combined action of many SN explosions, in a cosmological resimulation of a low mass galaxy at $z \ge 9$ from the `{\sc nut}' suite. However, in contrast with previous work, we achieve a resolution high enough to capture {\em individual} SN remnants in the Taylor-Sedov phase, for which the Taylor-Sedov solution actually provides an accurate description of the expansion. We report the development of a high-velocity, far-reaching galactic wind produced by the combined action of SNe in the main galaxy and its satellites, which are located in the same or a neighbouring dark matter halo. Despite this, we find that (i) this wind carries out very little mass (the measured outflow is of the order of a tenth of the inflow/star formation rate) and (ii) the cold gas inflow rate remains essentially unchanged from the run without SNe feedback. Moreover, there are epochs during which star formation is enhanced in the feedback run relative to its radiative cooling only counterpart. We attribute this `positive' feedback to the metal enrichment that is present only in the former. We conclude that at very high redshift, efficient SNe feedback can drive large-scale galactic winds but does not prevent massive cold gas inflow from fuelling galaxies, resulting in long-lived episodes of intense star formation.

Significant outflows of gas from star-forming galaxies, known as galactic winds, have been observed both in local galaxies \citep[e.g.][]{lehnert_heckman_1996, cmartin99_outflows} and high redshift ($z \sim3$) Lyman break galaxies (LBGs) \citep[e.g.][]{pettini_etal_2001, shapley_etal_2003}. These winds are usually associated with starbursting galaxies (for example LBGs have star formation rates (SFRs) of $\sim 100 {\rm M}_{\odot}/{\rm yr}$) and in this case are often referred to as `superwinds'. These winds can reach large distances from the source galaxy and are typically observed to have temperatures greater than the `escape temperature' of the potential well of their host halo e.g. M82, a classic example of a superwind \citep[][]{lehnert_heckman_weaver_1999}. This suggests that at least some of their metal enriched gas may be able to escape into the IGM. The production of these superwinds is attributed to simultaneous SNe explosions whose remnants can overlap producing a bubble of hot, low density gas that can `blowout' of the galaxy and become a wind \citep{mckee_ostriker_1977}. The main requirement for `blowout' is that the rate of SNe is high enough, such that the remnants overlap before they can cool radiatively \citep{heckman_armus_miley_1990, david_forman_jones_1990}. Observations of outflows from M82 are consistent with multiple SNe supplying energy to drive the wind \citep{heckman_armus_miley_1990}. Alternatively, it has been advocated that multiple coherent SNe, as seen in the superwind scenario, are not the only mechanism able to produce outflows from dwarf galaxies. Dwarf galaxies can lose a significant fraction of their mass on timescales of a Gyr, in a more quiescent manner. Fuelled only by an average SFR, cold gas clouds in the interstellar medium (ISM) are evaporated by SNe remnants and form a wind if the evaporated gas reaches high enough temperatures \citep{efstathiou_2000}. Due to their ability to eject material from galaxies, galactic winds are considered a potential solution to several outstanding problems in the field of structure formation. They are proposed to play a role in the pollution of the IGM and intracluster medium with metals, the occurrence of metallicity gradients within individual galaxies \citep[e.g.][]{heckman_etal_2000}, the properties of dwarf galaxies \citep[][]{dekel_silk_1986} and even to provide a reservoir of hot gas for re-accretion by galaxies at lower redshifts \citep{oppenheimer_etal_2010}. The most influential role of galactic winds, however, is arguably their impact on the IGM. \citet{furlanetto_loeb_2003} demonstrate analytically that LBGs can propagate $\sim 100$kpc into the IGM, which is compatible with the sizes of observed HI-deficient regions around LBGs \citep{adelberger_etal_2003}. Based on cosmological simulations, \citet{aguirre_etal_2001} argue that SNe driven winds from massive galaxies alone ($M_{\rm baryon}>10^{10.5} {\rm M}_{\odot}$) could pollute the entire IGM to its observed level. It has also been proposed, however, that protogalaxies (with masses comparable to present day dwarf galaxies) could be better candidates for polluting the IGM, since the lower velocity winds they would produce at $z=9$ would not perturb the IGM in the way that superwinds at $z=3$ would \citep[][]{madau_ferrara_rees_2001}. There is also evidence for metals in the IGM by $z=5$, requiring pollution at an earlier epoch than that probed by observations of LBGs. Simulations of an isolated dwarf by \citet{maclow_ferrara_1999} show that almost all metals are ejected from dwarf galaxies (with $M_{\rm gas}<10^9 {\rm M}_{\odot}$), even though the ejection of gas mass is surprisingly inefficient. It is clear that galactic winds have the potential to impact vast regions of space and influence structure formation on large scales. Therefore, the importance of studying this phenomenon in its full cosmological context cannot be overestimated. However, while there have been many successful numerical simulation studies of galactic winds in individual galaxies \citep{maclow_ferrara_1999,mori_ferrara_madau_2002, scannapieco_bruggen_2010}, capturing the production of a galactic wind in cosmological simulations remains a challenge. In such simulations, galactic winds tend to be imposed directly, rather than allowed to develop naturally as a result of overlapping SNe bubbles \citep{springelhernquist03_sf, scannapieco_etal_2006,oppenheimer_dave_08}. We recall that an important ingredient of analytical studies of SNe driven winds is a multiphase ISM i.e. a hot, rarefied atmosphere, containing cold dense gas clouds \citep[e.g.][]{mckee_ostriker_1977,efstathiou_2000}. If the resolution of a cosmological simulation is not sufficient, the multiphase nature of the ISM is not captured and it is not possible to accurately model SNe remnants. In fact, the inclusion of a multiphase ISM `by hand' is a common feature of galactic wind models \citep[e.g.][]{springelhernquist03_sf, scannapieco_etal_2006}. \begin{figure*} \includegraphics[width=0.31\textwidth,trim = 0mm 0mm 0mm 0mm, clip]{allgridswholebox.eps} \includegraphics[width=0.31\textwidth,trim = 0mm 0mm 0mm 0mm, clip]{allgrids6rvir.eps} \includegraphics[width=0.31\textwidth,trim = 0mm 0mm 0mm 0mm, clip]{allgridsrvir.eps}\\ \includegraphics[width=0.31\textwidth,trim = 0mm 0mm 0mm 0mm, clip]{densmap_fb9_293_level9_wholebox.eps} \includegraphics[width=0.31\textwidth,trim = 0mm 0mm 0mm 0mm, clip]{densmap_lmax14_00293_FB9_6rvir_z_subs_nocb.eps} \includegraphics[width=0.31\textwidth,trim = 0mm 0mm 0mm 0mm, clip]{densmap_lmax16_00293_FB9_rvir_z_subs_nocb.eps} \caption{Maps showing a z-projection of the hierarchy of grids (first row) and corresponding gas density (second row) in the feedback run at $z=9$. From left to right: The whole box, a cube of side $12r_{\rm vir} \approx 60$kpc centred on the main halo and a cube of side $2r_{\rm vir}$ also centred on the main halo respectively. The white circles in the middle and right gas density images (second row) indicate the virial radius of the main halo and its subhaloes. In the top left panel showing the whole box, we can see the $128^{3}$ grid covering the whole image and the 3 nested grids of equivalent resolution $256^{3}$ (orange), $512^{3}$ (yellow) and $1024^{3}$ (green) centred on our main halo. Additional grids can also be seen (blue,pink) indicating where additional levels of AMR have been triggered in the $1024^{3}$ grid. The middle and right panels show the triggering of additional levels of AMR (up to a maximum of 15) in dense structures, such as the main galaxy disc and subhaloes.} \label{grids} \end{figure*} Examining galactic winds in a cosmological context, rather than in isolation, introduces further complications. Galaxies are now embedded in the cosmic web and are subject to inflows as well as outflows. The mode of inflow will be dependent on the galaxy halo mass; low-mass galaxy haloes ($M_{\rm vir} < M_{\rm shock} \approx 10^{12} M_\odot$ ) cannot host a stable shock at the virial radius and so primarily accrete cold ($T < T_{\rm vir}$) gas whereas high-mass galaxy haloes ($M_{\rm vir} > M_{\rm shock}$) primarily accrete shock-heated, hot ($T \sim T_{\rm vir}$) gas \citep{shock_1d}. It has been suggested that this framework, when coupled with active galactic nuclei (AGN) and SNe feedback, could naturally give rise to some of the bimodality in galaxy properties we observe but usually fail to reproduce with simulations \citep{mshock}. Indeed, semi-analytic modelling has shown that incorporating such processes results in an improved fit to most galaxy property trends e.g. colour bimodality, luminosity function etc \citep{cattaneo}. \citet{keres2b} have shown, however, that AGN radio-mode feedback (which can prevent hot gas from being accreted) has little effect on high-mass galaxies in cosmological simulations because most of these were built hierarchically from lower mass galaxies which gained their baryons via cold accretion. They suggest it is feedback in low mass galaxies, such as SNe driven winds, that is of greatest importance. It seems that studying accretion and feedback in low mass protogalaxies is, therefore, vital for our understanding of galaxy evolution as a whole. Recent cosmological simulations are in agreement that for low mass galaxies (with $M_{\rm vir} < M_{\rm shock}$) most inflowing gas is cold \citep[e.g.][]{keres,bimodal_marenostrum,brooks}. All of these studies also demonstrate qualitatively that much of this cold inflow occurs along filaments (e.g. Fig.~17 of \citet{keres}, Fig.~5 of \citet{bimodal_marenostrum} and Fig.~5 of \citet{brooks}). What has not been quantified, however, is how much cold accretion occurs in a spherically symmetric way i.e. with the same geometry as assumed for hot accretion. While it is reasonable to assume that this spherical cold accretion can be safely ignored when the galaxy is surrounded by a halo of shock-heated gas (i.e. it has $M_{\rm vir} \sim M_{\rm shock}$), this is not the case when the halo is below the shock mass threshold and there is not necessarily anything to impede the spherical accretion of cold material from the IGM. Indeed, most of the visual demonstrations that cold accretion takes place along filaments in these previous studies focus on examples when the galaxy has a hot gas halo and cold accretion is necessarily confined to the filaments anyway. This paper thus complements such studies by measuring how cold gas accretion is divided between the filamentary and diffuse (spherically symmetric) components in low mass ($M_{\rm vir} < M_{\rm shock}$) haloes, in simulations both with and without SNe feedback. Most of the aforementioned studies also incorporate SNe feedback, although it is unclear whether a substantial galactic wind ever develops and they do not address the specific impact of this on the accretion processes. Recent work by \citet{vandevoort_etal_2010} attempts to tackle this issue with a suite of cosmological simulations employing different physics. We note, however, that their SNe feedback is implemented by giving gas particles a velocity kick and so the galactic wind is an input rather than a result that arises naturally. They conclude that the relative importance of cold accretion onto the galaxy halo is robust to changes to the feedback recipes, but that the relative importance of cold accretion onto the galaxy itself is decreased in the presence of SNe-driven winds. We revisit this question in this paper, but at much higher resolution and redshift ($z \ge 9$), where SNe blastwaves are individually resolved and so it is not necessary to put in a galactic wind by hand. We also emphasize that our high spatial and mass resolution allow us to fully resolve the filaments at very high redshift ($z \ge 9$), contrary to previous studies in which the spatial resolution, in particular, is lacking in this redshift range. These resolution effects potentially enhance the likelihood of the filaments being destroyed by the SNe feedback, changing the balance between inflow and outflow and thus the subsequent evolution of the galaxy. A few high resolution studies of cosmological accretion in individual galaxies have been performed, but these tend to focus either on more massive galaxies (i.e.~ with $M_{\rm vir} > M_{\rm shock}$) at lower redshift ($z \sim 3$) \citep{agertz_clumpygal,ceverino_clumpygal} or on the early formation stages of the first galaxies ($M_{\rm vir} \approx 5 \times 10^{7} M_\odot$) at $z>10$ where direct accretion of gas from the IGM briefly precedes cold filamentary inflows \citep{greif_etal_2008}. We are primarily interested in studying the epoch bracketed by these studies; that in which filamentary accretion is important for low mass haloes, during which galactic winds may also develop. In summary, for a low mass protogalaxy (i.e. with $M_{\rm vir} < M_{\rm shock}$) at high redshift ($\approx z=9-10$), there are two main predictions about its interaction with the cosmic web : 1) a galactic wind will develop and extend far into the IGM, polluting it with metals and 2) cold gas will flow rapidly into the host halo via the web's filaments. What happens when these inflow and outflow mechanisms occur simultaneously? We tackle this question with a suite of cosmological resimulations (incorporating different physics) with $0.5$pc physical resolution in the densest regions, of a galaxy with $M_{\rm vir} \sim 5 \times 10^{9} M_\odot$ at $z=9$. At such high resolution, we can model individual SNe with a Sedov blastwave allowing a galactic wind to arise naturally. We examine the mechanism via which the wind develops and measure its properties. Furthermore, by measuring the inflow and outflow rates and comparing these with a control run (without SNe) we investigate whether the presence of a hot galactic wind can alter the accretion processes in a protogalaxy in such a way as to significantly impact its evolution. \begin{figure*} \centering \includegraphics[width=0.4\textwidth,trim = 0mm 0mm 0mm 0mm, clip]{tempmap_lmax14_00101_COOL9_6rvir_z_subs.eps} \includegraphics[width=0.4\textwidth,trim = 0mm 0mm 0mm 0mm, clip]{tempmap_lmax14_00293_FB9_6rvir_z_subs.eps}\\ \includegraphics[width=0.4\textwidth,trim = 0mm 0mm 0mm 0mm, clip]{densmap_lmax20_00101_COOL9_10pcrvir_z.eps} \includegraphics[width=0.4\textwidth,trim = 0mm 0mm 0mm 0mm, clip]{densmap_lmax20_00293_FB9_10pcrvir_z.eps} \caption{{\bf Top row: }Density weighted temperature for a cube of $12 r_{\rm vir} \approx 60$ kpc on a side centred on the main halo for the cooling (left) and feedback (right) runs at $z=9$. Black circles in the top images indicate the position and virial radius of the main halo (the largest circle) and its subhaloes. The colour scale has been fixed to $3.0 \le \log T \le 5.3$; $\log T=5.3$ is the threshold for gas to be considered `hot' (see Table \ref{catsummary}) and the lower limit of $\log T=3$ has been chosen such that the filaments can be seen clearly since a significant mass fraction of the gas is in the range $3.3 \le \log T \le 4.3$. {\bf Bottom row: } Projected density for a cube of $0.2 r_{\rm vir}$ on a side centred on the main halo, showing the disc approximately face-on in the cooling (left) and feedback (right) runs at $z=9$. Note the clumps visible in the feedback run (right).} \label{tempmaps} \end{figure*} This paper is structured as follows. In Section \ref{sec:sim} we describe the suite of simulations we have performed and outline the gas physics employed (cooling, star formation, SNe feedback), further details of which are provided in the Appendix (Section \ref{sec:appendix}). In Section \ref{sec:filwind} we distinguish between different gas phases in both runs, relating these to physical structures e.g. the filaments, satellite galaxies, the hot galactic wind (feedback run only). In Section \ref{sec:hotgas}, we explore the mechanism responsible for the development of the galactic wind in the feedback run and measure the wind properties. We measure the inflow and outflow in all gas phases in Section \ref{sec:acc} and then relate the dominant accretion modes and outflows to the star formation rate in Section \ref{sec:sfr}, comparing the cooling and feedback runs in order to isolate the impact of the SNe feedback on the formation and evolution of this galaxy at high-redshift. Finally, we outline our conclusions in Section \ref{sec:conc}.

\label{sec:conc} We have undertaken a detailed analysis of a suite of subparsec resolution simulations (the {\sc nut} simulations) with the aim of understanding the mechanism via which a far-reaching galactic wind can arise in a protogalaxy at high redshift ($z\approx9$) and how this impacts the surrounding IGM and the galaxy's evolution. In particular we examine the relationship between cold-mode, filamentary accretion and star formation and explore the impact that the introduction of SNe feedback and the resulting wind has on these processes. \\ \noindent Our main findings are: \begin{enumerate} \item A far-reaching SNe driven wind (illustrated in Fig.~\ref{tempmaps}, top right) does develop; this is particularly significant since we do not directly model the wind in our feedback scheme, but instead we resolve the Sedov blastwaves of individual SNe and the wind arises naturally. We find that SNe explosions occur not only in the main progenitor, but also in subhaloes and neighbouring dark matter haloes (demonstrated in the temperature-radius histograms, Fig.~\ref{radhisto_fb9}). These create overlapping bubbles which are critical for creating the far-reaching wind and for facilitating the escape of the gas from the halo potential. Since subhaloes are key, we speculate that similar winds can also be produced by higher mass haloes at lower redshifts, and we will certainly investigate this issue in the very near future. The wind is enriched with metals to the order of $0.1 Z_{\odot}$ (see the metallicity profile in Fig.\ref{mwvrvr_big}, top right). While one must be careful when comparing high-redshift simulations results to local observations, we conclude that it seems feasible for such a wind to enrich the IGM to the level required by observations at $z \approx 2$ . On the other hand, we find the ejection of mass by the wind is very inefficient: the outflow rate in the hot phase is only $10-30$ per cent of the total mass inflow rate. We also find that the mass outflow rate is about an order of magnitude lower than the SFR, in disagreement with local observations \citep[e.g][]{cmartin99_outflows, heckman_etal_2000}, perhaps indicating an important difference between low and high redshift SNe driven winds or a missing physical process in the simulation which drastically alters the mass loading factor of the wind (stellar winds, radiative transfer). We plan to address this issue in a forthcoming paper. \item The total gas accretion rate at the virial radius of the main progenitor (see Fig.~\ref{fluxoutvr_fb9}, top two rows) ranges between $1-10 {\rm M}_{\odot} {\rm yr}^{-1}$ depending on redshift, and is in fair agreement with analytical predictions based on extended Press-Schechter theory, indicating that the gas accretes like dark matter at the virial radius (note that the accretion is cold). However, in contrast to the accretion of dark matter, this rate persists for the gas filaments right down to the central object and thus these filaments are the dominant supply of cold gas to the central disc (compared to clumpy or diffuse, spherical cold accretion). While the filaments appear more perturbed in the inner halo in our 3D visualisations of the feedback run than in those of the run without SNe feedback (Fig.~\ref{3dfluxcuts}, second row), we measure similar mass accretion rates in both runs, suggesting that the majority of the gas survives its journey to the central object intact and gets replenished by cold gas coming in from larger distances. So, despite the development of a galactic wind, the cold, filamentary accretion is not significantly altered. This suggests that, at least at $z \approx 9$, SNe driven winds cannot reduce the amount of cold, accretion onto dwarf-mass protogalaxies and therefore cannot significantly reduce star formation at this epoch, as is commonly assumed. \item At the lower end of the redshift range studied ($z<12$), the SFR in the feedback run is more often than not greater than that in the cooling run (see Fig.~\ref{ratio}, left for the ratio of SFRs versus redshift). An extreme example is at $z=9$ when the SFR in the feedback run is $\approx 3 {\rm M_{\odot} {\rm yr}^{-1}}$, yet only $0.3 {\rm M_{\odot} {\rm yr}^{-1}}$ (the minimum SFR over the redshift range considered) in the cooling run. We attribute this `positive' feedback effect to the metal enrichment active only in the former (we have already demonstrated that the gas accretion rates are very similar in both runs). Increased metallicity means more efficient cooling and this combined with the wind which can spread the metals, means a greater mass of gas will cool and condense enough to form stars. In fact, we find that the mass outflow rate is significantly lower than the rate of input of ejecta and entrained ISM into the SNe blastwaves, suggesting much of this material is recycled for next generation star formation. \item Star formation is globally very efficient in both the cooling and feedback runs. Even though in the star formation implementation we set the efficiency to $0.01$ when the gas density in a cell goes over $10^{5}$ atoms/cc, we measure average efficiencies of $\sim 0.2$ within $0.1r_{\rm vir}$. This indicates that star formation at high redshift proceeds in dense clumps that represent a larger fraction of the total gas mass of the galaxy than they do at low redshift. The global efficiency is well correlated with the average inflow rate of clumpy gas (i.e. gas that is condensed in satellites/subhaloes) (see Fig.~\ref{ks}, bottom) i.e. with the occurrence of mergers or strong interactions between neighbouring galaxies. When mergers occur the ISM is compressed and/or turbulence triggers increased fragmentation, allowing a larger fraction of the gas to reach high enough densities for star formation, resulting in higher global star formation efficiencies. \end{enumerate}

1012.2839

1012

1012.1698_arXiv.txt

{Studying the composition of dust in the interstellar medium (ISM) is crucial in understanding the cycle of dust in our galaxy.} {The mid-infrared spectral signature of {amorphous silicates, the most abundant dust species in the ISM,} is studied in different lines-of-sight through the {Galactic} plane, thus probing different conditions in the ISM.} {We have analysed 10 spectra from the Spitzer archive, of which 6 lines-of-sight probe diffuse interstellar medium material and 4 probe molecular cloud material. The 9.7 $\mu$m silicate absorption features in 7 of these spectra were studied in terms of their shape and strength. In addition, the shape of the 18 $\mu$m silicate absorption features in 4 of the diffuse sightline spectra were analysed.} {The 9.7 $\mu$m silicate absorption {bands} in the diffuse sightlines show a strikingly similar band shape. This is also the case for all but one of the 18 $\mu$m silicate absorption bands observed in diffuse lines-of-sight. The 9.7 $\mu$m {bands} in the 4 molecular sightlines show small variations in shape. These modest variations in the band shape are inconsistent with the interpretation of the large variations in $\tau_{9.7}$/$E$(J$-$K) between diffuse and molecular sightlines in terms of silicate grain growth. Instead, we suggest that the large changes in $\tau_{9.7}$/$E$(J$-$K) must be due to changes in $E$(J$-$K).}{}

The composition of dust in the interstellar medium (ISM) is a result of {a variety of} different processes. First of all, dust is formed in the circumstellar environments of evolved stars, for example asymptotic giant branch (AGB) stars. Depending on their stage of evolution these stars produce either oxygen-rich (e.g. silicates and oxides) or carbon-rich dust (e.g. amorphous carbon, {graphite and silicon carbide}). The {newly-formed} {stardust} enters the ISM through stellar winds or supernova explosions. This material is rapidly mixed with other gas and dust in the ISM, where it is processed by a number of mechanisms {such as} shock waves driven by supernova explosions, high energy radiation and high velocity collisions amongst grains. Finally, during the formation of a new star and planetary system the dust is further processed. Therefore, the composition of dust in the ISM is a reflection of the constant formation, mixing, processing and destruction of different dust species. Studying the composition of interstellar dust and its spatial variations throughout the galaxy is crucial to get a better understanding of the material from which, for instance, our solar system was made. A good way to study the composition of interstellar dust is by means of infrared spectroscopy. The light {from the bright background star is attenuated} by the interstellar dust in front of it. This extinction is wavelength dependent and is caused by scattering as well as absorption by the dust grains. If the intrinsic, unreddened spectrum of the background star is known, the dust extinction as a function of wavelength can be derived. \begin{table*} \hspace*{-0.5cm} \begin{tabular}{llllllll} \hline \\ target name & {other} & l & b & Program ID & spect.& diffuse (D) or & 9.7 or 18 $\mu$m\\ & {designations} & (degrees) & (degrees) & & type & molecular (M) & silicate feature\\ \hline\\ StRS 136 & {2MASS J17475608} & 0.040937 & $-$0.566892 & 3616 & B8--A9I$^{2}$ & D & 9.7\\ & {$-$2911439} & & & & & & \\ StRS 164 & {2MASS J18161876} & 14.213568 & $-$0.002043 & 3616 & B8--A9I$^{2}$ & D & 9.7\\ & {$-$1635468} & & & & & & \\ {SSTc2d\_J182835.8+002616} & {\ldots} & 30.712112 & 5.280682 & 139 $\&$ 172-179 & \ldots & M (Serpens) & 9.7\\ {$[$SVS76$]$ Ser 9}$^{1}$ & {2MASS J18294508} & 31.626616 & 5.423514 & 139 $\&$ 172-179 & \ldots & M (Serpens) & 9.7\\ & {+0118469} & & & & & & \\ StRS 354 & {IRAS 20273+3740} & 76.972032 & $-$0.635686 & 20294 & O7--B3$^{2}$ & D & 9.7 \& 18\\ {Elias 3-13}$^{3}$ & {IRAS 04303+2609} & 172.694191 & $-$14.498000 & 139 $\&$ 172-179 & K2III$^{3}$ & M (Taurus) & 9.7\\ SSTc2d\_J163346.2$-$242753 & {\ldots} & 354.158006 & 15.606477 & 139 $\&$ 172--179 & \ldots & M ($\rho$ Ophiuchi) & 9.7 \\ {G323.2103$-$00.3473}$^{4}$ & {2MASS J15285631} & 323.211041 & $-$0.347425 & 3616 & \ldots & D & 18\\ & {$-$5653045} & & & & & & \\ {G343.6142$-$00.1596}$^{4}$ & {2MASS J16585873} & 343.613830 & 0.160012 & 3616 & \ldots & D & 18\\ & {$-$4220543} & & & & & & \\ {G345.3650$-$00.4015}$^{4}$ & {2MASS J17070616} & 345.364763 & $-$0.401650 & 3616 & \ldots & D & 18\\ & {$-$4118140} & & & & & & \\ \hline\\ \end{tabular} \caption{Overview of the selected spectra from the Spitzer archive. Listed are the target name and the {Galactic} longitude {and latitude}. The Spitzer ID of the program that contains the observation and the spectral type of the source {is} listed and it is indicated if the observation was done in a diffuse (D) or molecular (M) sightline. In the case of molecular sightlines, the corresponding molecular cloud is listed. Finally, we list whether the 9.7 or 18 $\mu$m silicate feature was extracted from these spectra. $^{1}$\citet{1976AJ.....81..314S}, $^{2}$\citet{Rawlings00}, $^{3}$\citet{Elias78}, $^{4}$\citet{2003yCat.5114....0E}.} \label{tab:samplesel} \end{table*} The most abundant dust species in the ISM are amorphous silicates, which cause two prominent absorption features at about 9.7 and 18 $\mu$m \citep[e.g.~][]{1971Natur.233...72S,1974ApJ...193L..81R,1975A&A....45...77G,1980ApJ...242..965M,1989ESASP.290...79R,1998MNRAS.298..131B,Kemper04,Chiar06,Min07}. {In this article we will discuss both the 9.7 $\mu$m and 18 $\mu$m bands.} Interstellar silicate dust is mostly a mixture of amorphous silicates with an olivine (O/Si=4) or pyroxene (O/Si=3) composition \citep[e.g.][]{1974ApJ...192L..15D, 1975A&A....45...77G, Kemper04, Min07, Chiar07}. Although olivines {and} pyroxenes are by definition crystalline silicates, in this paper we will also use these names for amorphous silicates {with the same stoichiometric} composition. From previous studies we know that the composition of interstellar silicates can vary, depending on the location in the ISM. For example \citet{Demyk99,Demyk00,Demyk01} showed that the dust around protostars has a higher O/Si ratio than the {newly-formed} dust observed around AGB stars. \citet{Chiar07} compared the strength of the 9.7 $\mu$m silicate absorption feature with the near-infrared colour excess $E$(J$-$K) in both diffuse and molecular sightlines\footnote[1]{Diffuse and molecular sightlines are defined as lines-of-sight which pass mostly, but not exclusively, through diffuse or molecular cloud material, respectively.}. {The 9.7 $\mu$m silicate feature, $\tau_{9.7}$, is due to the Si-O stretching mode in the silicates, and its strength depends on the silicate optical depth, while the overall extinction in the visible through {near-infrared} is probably due to a combination of carbonaceous material \citep[e.g.~amorphous carbon;][]{Draine84,CVJ_10_dustSED} and iron, the latter either in metallic form \citep{KDW_02_composition} or as a cationic part of the silicate lattice, the so-called \emph{dirty silicates} \citep{JM_76_dust}}. For diffuse sightlines there is a tight linear correlation between $\tau_{9.7}$ and the visual extinction ($A_{\mathrm{V}}$) \citep{Roche84,Whittet03}. A similar correlation exists between $\tau_{9.7}$ and the near-infrared extinction represented by the J-K colour excess, $E$(J$-$K) \citep{Chiar07}. However, for the molecular sightlines the diffuse ISM correlation fails \citep{Whittet88_2,Chiar07}, indicating that something happens to the dust grains when they enter a molecular cloud. Molecular sightlines, in general, probe {greater interstellar densities} than diffuse sightlines and thus, the observed variations are most likely density related. An obvious explanation is grain growth via coagulation taking place in molecular clouds. Indeed, this causes a decrease in the strength of the 9.7 $\mu$m silicate feature (see Sect.~\ref{sect:models}), but grain growth also has an effect on the {near-infrared} extinction (i.e. $E$(J$-$K)). Moreover, {grain growth} alters the \emph{shape} of the 9.7 $\mu$m silicate feature drastically {\citep[e.g.~][]{BMD_01_processing}}. {Therefore, an accurate measurement of the band shape of the 9.7 $\mu$m silicate band in diffuse and molecular lines-of-sight can set interesting limits on the importance of grain growth as an explanation for the large changes in extinction detected so far.} In order to study the spectral signature of interstellar silicates in both diffuse and molecular sightlines, we have analysed 10 low-resolution spectra {with high S/N ratio,} taken with the Spitzer space telescope. {The high quality of the spectra allows an accurate comparison of band shapes along different lines-of-sight at a variety of Galactic longitudes}. Of these spectra, 7 were used to analyse the shape and strength of the 9.7 $\mu$m silicate absorption feature and 4 were used to analyse the 18 $\mu$m silicate absorption profile. {Only a few earlier studies have analysed the interstellar 18 $\mu$m band, and always in conjunction with the 9.7 $\mu$m band \citep[e.g.~][]{Demyk99,Chiar06}.} This paper is organised as follows: Section 2 describes the sample selection and the observations. In Sect.~3 the data analysis is described and the results are presented. {These results are compared with model calculations in Sect.~4, and in Sect.~5 we discuss the implications of our results for grain processing in molecular clouds. Finally, the conclusions are summarised in Sect.~6.}

In order to study the dust properties in different environments in the ISM, spectra were selected from the Spitzer archive of highly {reddened stars} in both diffuse and molecular sightlines. We have analysed the 9.7 and 18 $\mu$m silicate absorption features in these spectra separately. Our main results are: \begin{itemize} \item{For diffuse lines-of-sight the strength of the 9.7 $\mu$m silicate feature represented by the optical depth at about 9.7 $\mu$m ($\tau_{9.7}$) shows a tight linear correlation with the near infrared colour excess{, $E$(J$-$K)}. However, this correlation breaks down for molecular sightlines \citep{Chiar07}. The measurements for the spectra analysed in {this paper} are consistent with this finding.} \item{The shape of the 9.7 $\mu$m silicate feature is {very} similar for the three diffuse sightlines investigated in this study. Their shape is, within errors, also similar to that observed in the line-of-sight towards the {Galactic Centre} \citep{Kemper04}. We conclude that we only observe small variations {less than 0.15 in normalised $\tau$} in the 9.7 $\mu$m silicate band shape in diffuse lines-of-sight, and that the {Galactic Centre} line-of-sight is representative for the diffuse ISM in general. {The shape of the 18 $\mu$m silicate absorption feature is also, {within errors}, found to be invariable for the diffuse sightlines.} The situation is different for lines-of-sight passing through molecular clouds: we observe small variations on the short wavelength side of the 9.7 $\mu$m silicate band.} \item{Since both the shape of the 9.7 $\mu$m silicate feature and the relationship between $\tau_{9.7}$ and $E$(J$-$K) change for molecular sightlines {as} compared to diffuse sightlines, we conclude that the dust properties in molecular clouds are different from the diffuse ISM.} \item{{Based} on comparison with theoretically calculated profiles, we have ruled out grain growth beyond about 1 $\mu$m as being responsible for the changes in the 9.7 $\mu$m profile. The observations could be explained if the dust grains are more spherical in molecular clouds than in the diffuse ISM, but the underlying process is not understood. } \item{We propose that the flattening of the $\tau_{9.7}$ versus $E$(J$-$K) relationship is caused by an increase in $E$(J$-$K) in molecular clouds. The mechanism that causes this, however, is not yet clear. Further research is necessary to investigate whether it can be caused by grain growth.} \end{itemize}

1012.1698

1012

1012.3462_arXiv.txt

We constrain the linear and quadratic bias parameters from the \emph{configuration} dependence of the three-point correlation function (3PCF) in both redshift and projected space, utilizing measurements of spectroscopic galaxies in the Sloan Digital Sky Survey (SDSS) Main Galaxy Sample. We show that bright galaxies ($M_r < -21.5$) are biased tracers of mass, measured at a significance of $4.5\sigma$ in redshift space and $2.5\sigma$ in projected space by using a thorough error analysis in the quasi-linear regime ($9-27\; \hmpc$). Measurements on a fainter galaxy sample are consistent with an unbiased model. We demonstrate that a linear bias model appears sufficient to explain the galaxy-mass bias of our samples, although a model using both linear and quadratic terms results in a better fit. In contrast, the bias values obtained from the linear model appear in better agreement with the data by inspection of the relative bias, and yield implied values of $\sigma_8$ that are more consistent with current constraints. We investigate the covariance of the 3PCF, which itself is a measurement of galaxy clustering. We assess the accuracy of our error estimates by comparing results from mock galaxy catalogs to jackknife re-sampling methods. We identify significant differences in the structure of the covariance. However, the impact of these discrepancies appears to be mitigated by an eigenmode analysis that can account for the noisy, unresolved modes. Our results demonstrate that using this technique is sufficient to remove potential systematics even when using less-than-ideal methods to estimate errors.

\label{s:intro} Studying the statistical properties of the galaxy distribution allows one to probe the structure of overdense regions today, learning about galaxy formation and cosmology. We observe significant clumping in this large-scale structure (LSS), which is commonly characterized by a series of \npoint\ correlation functions \citep[reviewed in][]{peebles:80}. Observational evidence is in line with predictions of a dark-energy dominated cold dark matter ($\lcdm$) model \citep{komatsu:09,sanchez:09,reid:10}. However, there is a large conceptual hurdle between following the evolution of mass densities in gravitational collapse \citep[e.g. ][]{lss_review} and that realized by galaxy positions. A priori, there is little reason to believe a one-to-one correspondence exists between mass overdensities and galaxy positions; complex galaxy formation processes such as merging and feedback should have significant contributions. For example recent results from the Sloan Digital Sky Survey \citep[SDSS; ][]{york:00} in \citet{zehavi:05,zehavi:10} show clustering varies with galaxy luminosity and color. This discrepancy between the observed ``light'' in galaxies relative to the predicted ``mass'' clustering is often described as \emph{galaxy-mass bias}. The parameterization of galaxy-mass bias enables a two-pronged approach to probe both cosmology and galaxy formation. On one side, we map the clustering of galaxies to that of the underlying mass distribution allowing us to understand and constrain cosmology. Alternatively, the parameterization of the bias itself encodes useful information concerning galaxy formation processes. This approach distills observational data from hundreds of thousands of galaxies available in modern surveys, such as the the two-degree field galaxy redshift survey \citep[2dFGRS;][]{2dFGRS} and the SDSS into a significantly smaller and more manageable form. Most observational evidence exploits the two-point correlation function (2PCF), the first in the series of \npoint\ functions (or equivalently, the power spectrum in Fourier space). However, the 2PCF represents only a portion of the available information. Measurements of higher order moments, such as the three-point correlation function (3PCF), allow a more complete picture of the galaxy distribution. The statistical strength of higher order information might rival that of two-point statistics \citep{sefusatti:05}, as well as break model degeneracies describing cosmology and galaxy bias \citep{zheng:07,kulkarni:07}. Previous analyses have estimated the 3PCF from modern galaxy redshift surveys, including work on the the 2dFGRS \citep{jing:04,wang:04,gaztanaga:05} and results from SDSS data \citep{kayo:04,nichol:06,kulkarni:07,gaztanaga:09,marin:11}. Related higher order statistics have also been measured for these datasets \citep{verde:02,pan:05,hikage:05,nishimichi:07}. This work is the second of two papers analyzing the reduced 3PCF on SDSS galaxy samples. The first paper \citep{mcbride:10} focused on the details of the measurements we analyze here, as well as clustering differences due to galaxy luminosity and color. This paper utilizes the configuration dependence to constrain non-linear galaxy-mass bias parameters in the local bias model \citep{fry:93}, and the properties of the errors necessary for quantitative analyses. The local bias model is a simple approach to characterize galaxy-mass bias. Alternative descriptions exist based on the halo model \citep[reviewed in][]{cooray:02}, which form phenomenological models with a wider range of parameters. Two well used formulations include the halo occupation distribution \citep[HOD;][]{berlind:02} and the conditional luminosity function \citep[CLF;][]{yang:03,vdB:03}. There are formulations for the 3PCF; however, the accuracy of the model predictions is not as well determined as the 2PCF when compared with data \citep[see][]{takada:03,wang:04,fosalba:05}. A significant advantage of a HOD modeling is the ability to use well determined measurements of the small scales for constraints (the non-linear regime in gravitational perturbation theory). Understanding the projected 3PCF, $Q_{proj}$, a major component of this work, provides a critical link to obtain reliable measurements at these smaller scales from observational galaxy samples. However, by using this simple prescription for galaxy-mass bias, we investigate the effects of binning and covariance resolution in a quantitative analysis with a clear and simple model where the implications for bias and cosmology are better studied for higher order moments. An important part of our analysis is comparing results from the projected 3PCF with the more commonly used redshift space measurements. This paper is organized as follows. We discuss the SDSS data, simulations, and mock galaxy catalogs in \S\ref{s:data}. We review the theory and methods of our analysis in \S\ref{s:methods}. We constrain the non-linear galaxy mass bias parameters in \S\ref{s:bias}. In \S\ref{s:ev}, we investigate clustering properties contained in the eigenvectors of the 3PCF covariance matrix. We perform a detailed examination of the quality of error estimation in \S\ref{s:errors}. We discuss our results and compare to related analyses in \S\ref{s:disc}. Finally, we review our main conclusions in \S\ref{s:summary}. Unless otherwise specified, we assume a flat $\lcdm$ cosmology where $\Omegam = 0.3$, $\OmegaL = 0.7 $, and $H_o = h \, 100 \Hunits$, used to convert redshift to physical distances.

\label{s:disc} We utilize the the configuration dependence of the 3PCF in redshift and projected space to constrain galaxy-mass bias parameters in the local bias model. We find that galaxies are biased tracers of mass, with brighter galaxies corresponding to increased bias. These results are consistent with detailed analysis of SDSS galaxies from the 2PCF \citep{zehavi:05,zehavi:10} which quantifies how bias increases clustering for brighter galaxy samples. Our results indicate that a linear bias model yields reasonable approximations to the observations, in agreement with \citet{hikage:05}. However, a non-linear bias model produces slightly better agreement, and yields lower reduced chi-square values ($\chi^2_\nu$ in Tables~\ref{t:bias} and \ref{t:linbias}). We notice a strong correlation between linear and quadratic bias, as expected from inspection of \eqref{eq:biasQ}, and consistent with measurements of SDSS galaxies using the bispectrum \citep{nishimichi:07}. We find that our redshift space measurements predict significantly negative quadratic bias with a linear bias near one. This effect was seen in a similar analysis conducted on 2dFGRS galaxies \citep{gaztanaga:05}. Interestingly, we find projected measurements suggest a larger linear bias with near zero quadratic bias for the same samples, suggesting a possible systematic effect from redshift distortions in this simple bias model. We examined the relative bias in \S\ref{ss:brel}. We find supporting evidence that the brighter galaxy sample is a \emph{more biased} realization using both the 2PCF and 3PCF, consistent with other analyses of SDSS data \citep{zehavi:02,zehavi:05,zehavi:10}. Relative bias provides a consistency check on the ``absolute'' galaxy-mass bias parameters we constrain, suggesting a combination of linear and quadratic bias terms are consistent with observations. However, the relative bias of the 2PCF suggests that our two parameter bias model fits underpredict the value of linear bias necessary to explain the observations. Again, we see a hint that constraints from projected measurements appear to be less affected -- although we caution that this trend has weak statistical significance given the larger uncertainties in projected space. We obtain reasonable projections for $\sigma_8$ by using our linear bias values from fits on $Q(\theta)$ in conjunction with the 2PCF. We estimate the values of $\sigma_8$ to be between $0.83$ and $1.13$ based on the BRIGHT ($M_r < -21.5$) and LSTAR ($-21.5 < M_r < -20.5$) galaxy samples. The values we obtain are contingent on a specific model of mass clustering, where we have chosen to use \nbody\ simulations (specifically the Hubble Volume \lcdm\ results), and redshift distortions (which we include through velocity information to distort particle positions in the HV simulation). For comparison, constraints of $\sigma_8$ from a joint analysis of the cosmic microwave background (CMB), supernova data (SN) and baryon acoustic oscillations (BAO) find $\sigma_8 = 0.82$ \citep{komatsu:09}. Our lower values are in good agreement with these constraints. Our high end values appear too large, but our results are in reasonable agreement with an analysis of a related statistic, the monopole moment of the 3PCF, where they find best fit $\sigma_8$ values between $0.9$ and $1.07$ \citep[see Table~3 in][]{pan:05} using 2dFGRS galaxies \citep{2dFGRS}. Although the value of $\sigma_8$ we obtain is comparable with results from 2dFGRS, the specific bias values will not be, as the 2dF targets a different galaxy selection than our SDSS samples. If we underestimate the value of linear bias, effectively $B$ here, \eqref{eq:bias2pt_s8} shows that the implied value of $\sigma_8$ will be overestimated. This might explain the larger values of our estimates in comparison to WMAP analyses. Our projections for $\sigma_8$ use clustering measurements between $9-27 \; \hmpc$ and exploit only the configuration dependence of $Q(\theta)$. This is a much smaller slice of data that is significantly different than either the monopole measurement (which utilizes a larger range of scales without configuration dependence) or WMAP results (that combines a immense amount of data from both CMB and LSS analyses). We do not intend this analysis to complete with these constraints, but rather to help illuminate the role of galaxy-mass bias in future constraints of $\sigma_8$ using the 3PCF. Understanding the properties of measurement errors and the impact of empirical methods of estimating the covariance is a critical component necessary for quantitative constraints. Recent results have done comparisons on lower order statistics, such as the work by \citet{norberg:09}. We compared several properties of 3PCF covariance matrices estimated from jackknife re-sampling to those constructed from many realizations of independent galaxy mock catalogs. While we noted some concerning discrepancies, we found these typically affected only the least significant eigenmodes. We found many similarities between the covariance estimates, including physical descriptions for the first three eigenmodes which account for an overwhelming majority of the variance. We established the need to trim noisy, unresolved modes from the covariance. When trimmed, and the eigenmode analysis is properly utilized, we noted only a few significant differences, mostly in the case of $15$ jackknife samples. We conclude that our use of $30$ jackknife samples does not significantly affect our analysis. We analyze measurements of the configuration dependence of reduced 3PCF for two SDSS galaxy samples that were first presented in \citet{mcbride:10}. In both redshift and projected space, we characterize the galaxy clustering differences with those predicted by the non-linear mass evolution in the \lcdm\ Hubble Volume simulation. Here, we summarize our main results: \begin{itemize} \item We demonstrate that brighter galaxies remain a more biased tracer of the mass field by constraining the linear and quadratic galaxy-mass bias parameters using a maximum likelihood analysis on scales between $6$ and $27 \hmpc$. Conservatively using scales above $9 \hmpc$, the BRIGHT sample is biased at greater than $2\sigma$ and the fainter LSTAR shows no significant bias, in generally agreement with expectations from previous analyses of SDSS galaxies \citep{zehavi:05,zehavi:10}. The bias parameters and their significance are summarized in Table~\ref{t:bias}. \item We resolve the degeneracy between the linear and quadratic bias terms, which helps to explain the weak luminosity dependence observed in the reduced 3PCF. \item We find a linear bias model appears sufficient to explain the measurements of the 3PCF by re-fitting the linear bias while constraining the quadratic bias at zero (results reported in Table~\ref{t:linbias}). However, we find the two parameter fit is preferred in our likelihood analysis, as it yields a lower chi-square in the best fit value. \item The relative bias between samples of different luminosities (which is independent of the mass predictions), as well as the cosmological implications for values of $\sigma_8$ , show general consistency with previous analyses. Inspection of our results suggest that the linear bias values obtained without a quadratic bias term are preferred. This suggests that two-parameter bias constraints might underpredict the linear bias. \item We decompose the structure of the normalized covariance matrix as an alternative view into clustering properties of our samples. The eigenvectors of the first three dominant modes show coherent structure consistent with variations seen in the $Q(\theta)$ measurements, supporting our claim that the covariance is signal dominated and sufficiently resolved. \item We find that jackknife re-sampling methods cannot reproduce the correlation seen in the a 3PCF covariance matrix estimated from many realizations of mock galaxy catalogs. By performing a detailed comparison of the properties and structure of the errors, we identify that noisy, unresolved modes introduce significant discrepancies. We find that using an eigenmode analysis can mitigate the differences and conclude that our analysis should not be significantly affected by less-than-ideal methods of error estimation. \item Comparing results between redshift space and projected measurements implies a potential systematic bias on values from the redshift space analysis when scales below $9 \hmpc$ are included, which have been utilized in other comparable analyses. Since the small scale measurements contain more constraining power than larger scales, they drive the likelihood analysis even when larger scales are considered. \item On scales above $9 \hmpc$, the statistical significance of constraints from redshift space analyses appear stronger than those found in analyses of projected measurements. We attribute this result to the increased uncertainties of the projected 3PCF, which mixes in larger scales (with larger errors) due to the line-of-sight projection. When considered with the results of \citet{mcbride:10}, which finds the projected 3PCF recovers configuration dependence at small scales lost in redshift space, a combination of redshift space analysis at large scales and projected measurements at small scales would form a nice complement in future analyses. \end{itemize}

1012.3462

1012

1012.0846_arXiv.txt

We derive all single-field cosmologies with unit sound speed that generate scale invariant curvature perturbations on a dynamical attractor background. We identify three distinct phases: slow-roll inflation; a slowly contracting adiabatic ekpyrotic phase, described by a rapidly-varying equation of state; and a novel adiabatic ekpyrotic phase on a slowly expanding background. All of these yield identical power spectra. The degeneracy is broken at the 3-point level: unlike the nearly gaussian spectrum of slow-roll inflation, adiabatic ekpyrosis predicts large non-gaussianities on small scales.

Our starting point is the quadratic action for $\z$, assuming $c_s =1$: \be S = M_{\rm Pl}^2 \int \dtau \dcx z^2 \left\{ ( \z' )^2 - ( \grad \z )^2 \right\}, \ee where $z \equiv a \sqrt{2 \e}$ and primes denote derivatives with respect to conformal time $\tau$. This yields the mode function equation for the canonically-normalized variable $v = z\,\zeta$: \be v_k''+\left(k^2-\frac{z''}{z}\right) v_k = 0\,, \label{veom} \ee where $k$ is the comoving wavenumber. To generate a scale invariant spectrum from adiabatic initial conditions, it is sufficient for $z$ to satisfy \be \frac{z''}{z} = \frac{2}{\tau^2}\,. \label{z''} \ee Indeed, the solution to~(\ref{veom}) in this case is \be v_k = \frac{1}{\sqrt{2 k}M_{\rm Pl}} e^{-i k \t} \left(1-\frac{i}{k \t}\right)\,, \label{vk} \ee which implies that $k^{3/2} |\zeta_k| = \sqrt{ 1 + k^2 \tau^2 }/ \sqrt{2} M_{\rm Pl} z |\tau|$. As $\tau\rightarrow 0$, $k^{3/2} \abs{\z_k}$ is independent of $k$, as desired. In addition to generating a scale invariant $\z_k$, our background must be a dynamical attractor. Since $\z_k \sim 1/z |\tau|$ as $k\rightarrow 0$, the desired solution to~(\ref{z''}) is \be z \equiv \frac{\sqrt{2}}{m \ta}\,, \label{zed} \ee where $m$ is an arbitrary scale. Combining~\eqref{vk} and~\eqref{zed} yields $k^{3/2} \abs{\z_k} = m\sqrt{1 + k^2 \t^2}/2M_{\rm Pl}$, which is both scale invariant as $\t \gt 0$ and constant as $k \gt 0$. The observed amplitude of $\zeta\sim 10^{-5}$ fixes $m\sim 10^{-5}M_{\rm Pl}$. We pause to note that in an inflationary context the freeze-out or {\it $\z$-horizon} $\ta$ is usually identified with the {\it comoving Hubble horizon}, $h^{-1} \equiv 1/aH = a/a'$, but that more generally ({\it e.g.}, when $\e$ varies rapidly) the Hubble horizon and the $\z$-horizon can differ greatly. Using the definition $z = a \sqrt{2 \e}$,~(\ref{zed}) implies \be \e = \frac{1}{a^{2} m^{2} \t^{2}}\,. \label{ep} \ee Moreover, we can rewrite $\e = -\dot{H}/H^2 = {\rm d}H^{-1}/{\rm d}t$ in terms of the comoving Hubble horizon $h^{-1} = 1/aH$ as \be \left( h^{-1} + \t \right)' = \e\,. \label{hub} \ee Combining~\eqref{ep} and~\eqref{hub} then gives a second-order differential equation for $a(\tau)$. Instead, we will cast these as a pair of coupled first-order equations. By differentiating~\eqref{ep}, \be (\log \sqrt{\e})' = - \t^{-1} - h\,. \label{boxep} \ee Once we specify the signs of $h$ and $\t$,~\eqref{hub} and~\eqref{boxep} become coupled ODEs for $\h$ and $\e$. The behavior of~\eqref{boxep} will depend strongly on the relative magnitude of the Hubble horizon $\h$ and the $\z$-horizon $\ta$. We will therefore say that the Hubble horizon is {\it inside the $\z$-horizon} when $\h < \ta$, and {\it outside the $\z$-horizon} when $\h > \ta$. To solve these coupled equations, $h_{\rm fid}$ and $\epsilon_{\rm fid}$ must be specified at some fiducial time $\tau_{\rm fid} < 0$. To obtain a solution for $a(\tau)$, we can set $a_{\rm fid} = 1$ by a spatial rescaling $a\rightarrow \lambda a$, $\tau \rightarrow \tau/\lambda$. The equation of state is of course invariant, so $\epsilon_{\rm fid}$ fixes $\tau_{\rm fid}$ through~(\ref{ep}). In practice, we will specify not $|h_{\rm fid}^{-1}|$ but the ratio $|h_{\rm fid}^{-1}|/|\tau_{\rm fid}|$, which is invariant under the above spatial rescaling. \begin{figure} \centering \includegraphics[width=0.4\textwidth]{cases.pdf} \caption{Sketch of $|h^{-1}|$ for the contracting (dotted), expanding (dashed) and apex (thick dashed) branches of solutions.} \label{sketchsolns} \end{figure}

We have uncovered three distinct cosmological phases that yield a broad range of scale invariant modes: inflationary expansion, adiabatic ekpyrotic contraction~\cite{adiabaticek}, and adiabatic ekpyrotic expansion~\cite{austin}. All three phases generate identical power spectra for $\zeta$, and each is an attractor background. The degeneracy is broken at the 3-point level. The rapidly-varying equation of state characteristic of adiabatic ekpyrotic phases results in strongly scale-dependent non-gaussianities~\cite{adiabaticek2}. Our results imply that inflation is the unique single-field mechanism with unit sound speed capable of generating a broad range of scale invariant and gaussian modes. Forthcoming work~\cite{austin} will extend the analysis to include a general sound speed $c_s(\tau)$, the other degree of freedom at our disposal~\cite{piazza}. {\it Acknowledgments:} We thank D.~Baumann, A.~Joyce, L.~Leblond, J.-L.~Lehners, S.~Shandera, P.~J.~Steinhardt and M.~Zaldarriaga for helpful discussions. This work is supported in part by the US Department of Energy (DE-AC02-76-ER-03071) and the Alfred P. Sloan Foundation.

1012.0846

1012

1012.5652_arXiv.txt

We present the results from 10 nights of observations of the globular cluster NGC~6981 (M72) in the $V$, $R$ and $I$ Johnson wavebands. We employed the technique of difference image analysis to perform precision differential photometry on the time-series images, which enabled us to carry out a census of the under-studied variable star population of the cluster. We show that 20 suspected variables in the literature are actually non-variable, and we confirm the variable nature of another 29 variables while refining their ephemerides. We also detect 11 new RR Lyrae variables and 3 new SX Phe variables, bringing the total confirmed variable star count in NGC~6981 to 43. We performed Fourier decomposition of the light curves for a subset of RR Lyrae stars and used the Fourier parameters to estimate the fundamental physical parameters of the stars using relations available in the literature. Mean values of these physical parameters have allowed us to estimate the physical parameters of the parent cluster. We derive a metallicity of [Fe/H]$_{\mbox{\scriptsize ZW}} \approx -$1.48$\pm$0.03 on the \citet{zin1984} scale (or [Fe/H]$_{\mbox{\scriptsize UVES}} \approx -$1.38$\pm$0.03 on the new \citet{car2009} scale) for NGC~6981, and distances of $\sim$16.73$\pm$0.36~kpc and $\sim$16.68$\pm$0.36~kpc from analysis of the RR0 and RR1 stars separately. We also confirm the Oosterhoff type I classification for the cluster, and show that our colour-magnitude data is consistent with the age of $\sim$12.75$\pm$0.75~Gyr derived by \citet{dot2010}.

\label{sec:introduction} The study of Galactic globular clusters is important for many reasons. These stellar systems represent some of the oldest, and consequently metal poor, stellar populations in the Galaxy, and their scrutiny allows us to glean information regarding the formation and early evolution of the Galaxy. The spatial distribution of the clusters reveals a different aspect of the Galactic structure than other stars in the Galaxy, and their orbits and tidal tails provide constraints on the Galactic potential. Of course, what we learn about globular clusters in our own Galaxy is also applicable to other galaxies as well. Globular clusters are also believed to be a close approximation to a stellar laboratory since a cluster's members were formed at the same time from the same primordial material with the same composition, leading to a homogeneity of certain properties within each cluster, but with differences in these properties between clusters. Although this paradigm is being challenged by the recent discovery in some globular clusters of multimodal main sequences and sub-giant branches (\citealt{pio2009}, and references therein), indicating the existence of multiple stellar populations, most globular clusters do not exhibit such obvious deviations from a simple stellar population and the paradigm still holds. There are $\sim$150 globular clusters in our Galaxy for which their fundamental properties, such as metallicity, distance, age and kinematics, have been estimated by various methods (\citealt{har1993}). One independent method for estimating at least some of these quantities is by studying the population of RR Lyrae variable stars present in most clusters. This method uses the fact that the light curve morphology of RR Lyrae stars is connected with their fundamental stellar parameters, and consequently quantities such as metallicity, absolute magnitude and effective temperature may be calculated from the fit parameters of the Fourier decomposition of their light curves using empirical, semi-empirical or theoretical relations published in recent years (\citealt{sim1993}; \citealt{jur1996}; \citealt{jur1998}; \citealt{kov1998}; \citealt{kov2001}; \citealt{mor2007}). Appropriate mean values of these fundamental parameters then enable similar estimates of the physical parameters of the parent cluster. As part of a series of papers on detecting and characterising the variable stars in globular clusters, and using the results to estimate the parameters of the host cluster (\citealt{are2004}; \citealt{laz2006}; \citealt{are2008a}; \citealt{are2008b}; \citealt{are2010}), we have performed CCD time-series photometry of the globular cluster NGC~6981 (RA~$\alpha = 20^{\mbox{\scriptsize h}} 53^{\mbox{\scriptsize m}} 27.9^{\mbox{\scriptsize s}}$, Dec.~$\delta = -12\degr 32\arcmin 13\arcsec$, J2000; $l = 35.16\degr$, $b = -32.68\degr$) using the method of difference image analysis (Section~\ref{sec:observations_reductions}). The known variables in this cluster, which are exclusively RR Lyrae variables, have been studied in a handful of photographic observing campaigns (\citealt{sha1920}; \citealt{ros1953}; \citealt{saw1953}; \citealt{nob1957}; \citealt{dic1972b}), the most recent of which is now 40 years in the past. Periods and ephemerides are poorly determined or non-existent for a substantial number of variables, and light curves for many of the claimed ``variables'' have not been published. Therefore, in Section~\ref{sec:variable_stars}, we use our precision differential photometry to perform an essential variable star census for the cluster. In Section~\ref{sec:RRL_physical}, we use the Fourier decompostion of the RR Lyrae star light curves to estimate their fundamental physical parameters, and then in Section~\ref{sec:cluster_properties}, we use the RR Lyrae properties to estimate the metallicity of, and distance to, NGC~6981. We also discuss the age estimates for the cluster that are available in the literature in the context of our colour-magnitude diagram. Our conclusions are presented in Section~\ref{sec:conclusions}. Throughout this paper we adopt the RR Lyrae nomenclature introduced by \citet{alc2000}; namely, RR0 designates an RR Lyrae star pulsating in the fundamental mode, and RR1 designates an RR Lyrae star pulsating in the first-overtone mode.

\label{sec:conclusions} We employed the technique of difference image analysis to perform precision differential photometry on a set of CCD time-series observations of the globular cluster NGC~6981 in the Johnson $V$, $R$ and $I$ filters. The difference imaging technique has allowed us to study in detail the variable star population of the cluster, including in the highly crowded central regions where traditional PSF fitting fails. Compared to previous photographic studies of this cluster, our photometry is deeper by $\sim$4~magnitudes, achieves a much better precision per data point ($\le$20~mmag down to $\sim$18.5~mag as opposed to $\sim$50~mmag at $\sim$17~mag), and covers a longer time-base of $\sim$5~years with substantially more data points ($\sim$100 compared to $\sim$20-60). Consequently, we have been able to perform a census of the variable star population in NGC~6981, and to clarify the status of already known and suspected variables in the cluster. We have shown that 20 of the 58 stars that are labelled as variables in the literature are actually non-variable, and we present the details of all 43 confirmed variables in NGC~6981 in Table~\ref{tab:variables}. We have discovered 14 of the 43 confirmed variables in NGC~6981 from our data. The new detections consist of 11 RR Lyrae variables and 3 SX Phe variables, although we were only able to derive reliable periods for 6 of the newly discovered RR Lyrae stars. For the 27 of the 29 previously known RR Lyrae variables that lie in our field of view, we calculated a new set of ephemerides that improve substantially on their previous values in terms of accuracy. The current set of 43 variables in NGC~6981 consists of 40 RR Lyrae stars and 3 SX Phe stars. Furthermore, the set of RR Lyrae stars is made up of 31 RR0 type variables, 5 RR1 type variables, and 4 RR Lyrae variables with an ambiguous subtype, although these variables are most likely to be of the RR0 type (V44, V45, V52 and V53). Based on the mean periods for the RR0 and RR1 variables, and the ratio of the number of RR1 to RR0 type variables, we confirm the Oosterhoff type I classification for NGC~6981. We report the detection of a strong Blazhko effect in 5 RR0 variables (V11, V23, V28, V31 and V32), and a smaller amplitude Blazhko effect in another 5 RR0 variables (V10, V14, V15, V36 and V49), which implies a $\sim$34~per cent lower limit on the rate of incidence of the Blazhko effect in RR0 stars in NGC~6981. We also find that the RR0 star V29 exhibits a secular period change of $\beta\approx-$1.38$\times$10$^{-8}$~d~d$^{-1}$, indicating that the star pulsation frequency is slowly increasing over time. Our analysis of the light curves of the SX Phe variables has allowed us to detect two pulsation frequencies in the variable V54, corresponding to the fundamental and first overtone radial oscillation modes. We did not detect any other frequencies apart from the dominant frequency for the other two SX Phe variables, and therefore we are unable to identify their modes of oscillation. We provide calibrated $V$ and $I$ photometry in the Johnson-Kron-Cousins photometric system, derived from the analysis of a set of standard stars in the field of the cluster, and instrumental $r$ photometry, for all of the 41 confirmed variables in NGC~6981 in our field of view. This data is available in electronic form (see Supporting Information) and the format is shown in Table~\ref{tab:vri_phot}. In order to collect the information required to observe these objects into one place, we calculate celestial coordinates for each variable (see Table~\ref{tab:astrom}) and provide detailed finding charts in Figure~\ref{fig:finding_charts}. Finding charts for the two RR Lyrae variables V27 and V35 that are not in our field of view may be found in \citet{dic1972a}. We performed a Fourier decomposition of the light curves for 21 RR0 and 5 RR1 variables with reliable period estimates and suitable phase coverage, and we report the corresponding Fourier coefficients in Table~\ref{tab:fourier_coeffs}. The Fourier parameters have been used to estimate the metallicity, absolute magnitude, log-luminosity, and effective temperature for each RR Lyrae star based on empirical, semi-empirical, and theoretical relations available in the literature. Assuming that the RR Lyrae stars in NGC~6981 are of the same composition, distance and age, then appropriate averages of the derived properties of the RR Lyrae stars may be calculated and employed as estimates of these properties for the parent cluster. Applying this method, we derive a metallicity of [Fe/H]$_{\mbox{\scriptsize ZW}} \approx -$1.48$\pm$0.03 on the ZW scale for NGC~6981, and [Fe/H]$_{\mbox{\scriptsize UVES}} \approx -$1.38$\pm$0.03 on the new \citet{car2009} scale. Similarly, we derive mean true distance moduli of $\mu_{0} \approx\,$16.117$\pm$0.047 and 16.111$\pm$0.047~mag for the RR0 and RR1 variables, respectively, and corresponding distances to NGC~6981 of $\sim$16.73$\pm$0.36~kpc and $\sim$16.68$\pm$0.36~kpc, respectively. The age of NGC~6981 has been estimated by \citet{dot2010} as 12.75$\pm$0.75~Gyr from analysis of {\it Hubble Space Telescope} ACS observations, and our calibrated $V-I$ CMD cannot compete with their data in terms of accuracy and lack of systematic errors. Hence we simply demonstrate that our CMD data is consistent with the isochrones of the \citet{van2006} stellar evolutionary models when interpolated to the age and metallicity of NGC~6981, and if one chooses to ignore the colour effect of the reddening. It is likely that the reddening is smaller than we assumed in Section~\ref{sec:distance} or that there is a systematic error in the colour of the theoretical stellar evolutionary models that we adopted, or both.

1012.5652

1012

1012.2244_arXiv.txt

{Using the technique of spectral disentangling, it is possible to determine the individual spectra of the components of a multiple star system from composite spectra observed at a range of orbital phases. This method has several advantages: it is unaffected by line blending, does not use template spectra, and returns individual component spectra with very high signal-to-noise ratios.} {The disentangled spectra of a binary star system are very well suited to spectroscopic analysis but for one problem: the absolute spectral line depths are unknown because this information is not contained in the original spectra (unless there is one taken in eclipse) without making assumptions about the spectral characteristics of the component stars. Here we present a method for obtaining the atmospheric parameters of the component stars by the constrained fitting of synthetic spectra to observed and disentangled spectra.} {Disentangled spectra are fitted using synthetic spectra and a genetic algorithm in order to determine the effective temperatures, surface gravities and relative light contributions of the two stars in a binary system. The method is demonstrated on synthetic spectra and then applied to the eclipsing binary V615\,Per, a member of the young open cluster NGC\,869 (h\,Persei).} {The method works well for disentangled spectra with signal-to-noise ratios of 100 or more. For V615\,Per we find a normal He abundance but an Mg abundance, which indicates bulk metallicity, a factor of two lower than typical for nearby OB stars.} {}

\label{sec:intro} The technique of {\em spectral disentangling} {(\sc spd)} allows the isolation of the individual spectra of the component stars of a double-lined spectroscopic binary system from a set of composite spectra observed over a range of orbital phases. It was originally formulated in the wavelength domain by \citet{SimonSturm94aa} and in the Fourier domain by \citet{Hadrava95aas}. The technique simultaneously returns the best-fitting individual spectra and the orbital velocity amplitudes of the two stars. A detailed overview of {\sc spd} can be found in \citet{PavlovskiHensberge09xxx}. Compared to other methods of radial velocity measurement, {\sc spd} has several advantages. Firstly, it is independent of template spectra so avoids any systematic errors due to spectral differences between the target and template stars. Secondly, it is not affected by the blending of spectral lines of the two stars \citep[see][]{Hensberge++00aa,MeClausen07aa}. Thirdly, the resulting disentangled spectra contain the combined signal of the input spectra \citep{PavlovskiMe09mn} so have a much higher signal-to-noise (S/N) ratio. There are two disadvantages of the {\sc spd} approach. The first of these is that the continuum normalisation of the input spectra has to be very good in order to avoid low-frequency spurious patterns in the resulting disentangled spectra \citep{Hensberge++08aa}. The second is that relative continuum light contributions of the two stars cannot be found using {\sc spd} as this information is itself not contained in the observed spectra, unless a spectrum has been obtained during an eclipse \citep{Ilijic+04aspc}. {\sc spd} is well suited to the spectral analysis of stars in binary systems. Each disentangled spectrum contains only features due to one star, so can be analysed using standard methods for single stars. The high S/N ratios of disentangled spectra are very helpful to this process, but the undetermined continuum light ratio between the component stars complicates the spectral analysis. In this work we present a method to fit synthetic spectra to disentangled spectra, where the atmospheric parameters of the stars are determined simultaneously with the relative light contributions of the stars. A genetic algorithm is used for the optimization in order to ensure that the best solution is found in a parameter space which suffers from strong degeneracies, in particular between effective temperature (\Teff) and surface gravity (\logg). An important application of {\sc spd} is the study of detached eclipsing binary star systems (dEBs). These represent the primary source of directly-measured masses and radii of stars, and as such are cornerstones of stellar physics \citep{Andersen91aarv,Torres++10aarv}. {\sc spd} can be used to measure the velocity amplitudes of the stars, which are necessary for the mass and radius measurements, simultaneously with obtaining the individual stellar spectra for spectral analysis \citep{PavlovskiHensberge05aa, Pavlovski+09mn, PavlovskiMe09mn}. A major advantage of dEBs to this process is that the surface gravities of the stars can be obtained to within $\pm$0.01\,dex from the mass and radius measurements: these parameters can then be fixed in the spectral analysis and thus the degeneracy between \Teff\ and \logg\ avoided \citep{Simon++94aa,Hensberge++00aa}. In this work we demonstrate the genetic algorithm approach to fitting disentangled spectra on the dEB V615\,Persei, a member of the young open cluster h\,Persei. \citet[][hereafter SMS04]{Me++04mn} obtained a series of high-resolution spectra of V615\,Per and analysed them with published light curves \citep{Krzesinski++99aa} to measure the masses (4.08 and 3.18 \Msun) to accuracies of 2\% and the radii (2.29 and 1.90 \Rsun) to 5\%, resulting in surface gravities measured to within 0.05\,dex. The \Teff\ values were found to be $15\,000 \pm 500$\,K and $11\,000 \pm 500$\,K. SMS04 found that stellar evolutionary models required a subsolar metal abundance ($Z \approx 0.01$) to reproduce the measured masses and radii of V615\,Per. The Perseus Double Cluster comprises h\,Persei (NGC\,869) and $\chi$\,Persei (NGC\,884). It has been extensively studied via deep CCD photometry \citep{Keller+01aj, MarcoBernabeu01aa, Slesnick+02apj, CapillaFabregat02aa, Currie+10apjs}, from which there is general agreement on its distance (2.3 to 2.4 kpc) and age (13--14\,Myr). But these studies assumed a solar chemical composition, and their results may be systematically wrong if this assumption is incorrect. Conflicting results on the chemical composition of the Perseus Double Cluster are present in the literature. Detailed abundance analyses based on high-resolution spectra of hot stars \citep{Lennon++88aa, Dufton+90aa} have challenged previous findings of low helium abundances \citep{Nissen76aa, KlochkovaPanchuk87sval, WolffHeasley85apj}. \citet{Dufton+90aa} and \citet{SmarttRolleston97apj} found an approximately solar metal abundance from high-resolution spectra, but this was not supported by \citet{Vrancken+00aa}. In this work we attempt to shed additional light on this subject by measuring the helium and metal abundances of the stars in the dEB V615\,Per.

Spectral disentangling is a method for obtaining the individual spectra of the components of a binary star system from composite spectra obtained at a range of orbital phases. A disadvantage of this method is that the continuum light ratios of the stars are not found, because this information is not present in the observed spectra without making assumptions about the spectral characteristics of the stars. We present the {\sc genfitt} program, which uses a genetic algorithm to fit synthetic spectra to the disentangled spectra of both components of a binary system simultaneously. It returns the best-fitting atmospheric parameters (\Teff\ and \logg) and the light contributions of the two stars. From tests with synthesized spectra we find that {\sc genfitt} performs extremely well in determining the light ratio. It also returns reliable \Teff\ and \logg\ values in those cases where $S/N$ of the input disentangled spectra is $\geqslant$$100$, which is the usual situation for observational studies. The light contributions of the two stars will normally sum to unity, which provides a useful constraint for {\sc genfitt}. Contaminating light from a third star can in principle be found, in cases when the light contributions of the two stars in the binary sum to less than unity. Once the light contributions of the stars have been found, their disentangled spectra can be renormalised to the correct continuum levels. The resulting spectra can then be analysed using standard methods for single stars. If the two stars are eclipsing, their surface gravity values may be found to high precision and accuracy from analysis of the orbital velocity amplitudes found by spectral disentangling and light curves covering the eclipses. As a demonstration of the method we applied {\sc genfitt} to spectra of the eclipsing system V615\,Per, a member of the h\,Persei open cluster. The metal abundance of this cluster is controversial (see Sect.\,\ref{sec:intro}) but important in measuring its distance by the isochrone method. The spectra were disentangled and fed into {\sc genfitt}, and an abundance analysis was performed on the resulting renormalised spectra. The atmospheric parameters returned by {\sc genfitt} are in good agreement with previous work (SMS04) but are more precise. We find a normal solar helium abundance for V615\,Per\,A (star\,B is cooler and has only weak helium lines). The magnesium abundances for both stars are lower than those found for nearby OB stars, indicating that h\,Persei has a subsolar metallicity. This is in agreement with the results of SMS04, based on the complimentary method of comparing the masses and radii of the stars to the predictions of theoretical stellar evolutionary models.

1012.2244

1012

1012.5839_arXiv.txt

In the recent paper of \citet{Hooper10} it was reported that $\gamma$-ray emission from the Galactic Center region contains an excess compared to the contributions from the large-scale diffuse emission and known point sources. This excess was argued to be consistent with a signal from annihilation of Dark Matter with a power law density profile. We reanalyze the \fermi\ data and find instead that it is consistent with the ``standard model'' of diffuse emission and of known point sources. The main reason for the discrepancy with the interpretation of \citet{Hooper10} is different (as compared to the previous works) spectrum of the point source at the Galactic Center assumed by \citet{Hooper10}. We discuss possible reasons for such an interpretation.

The origin of the emission from the Galactic Center (GC) at keV--TeV energies has been extensively discussed in the literature over last few years. In their recent paper, \citet{Hooper10} claimed that the $\gamma$-ray emission from the Galactic Center region, measured with the \fermi\ LAT instrument~\citep{FermiOverview} cannot be described by a combination of spectra of known point sources, diffuse emission from the Galactic plane and diffuse spherically symmetric component (changing on the scales much larger than $1^\circ$). An additional spherically symmetric component was suggested to be needed in the central several degrees. This component was then interpreted as a dark matter annihilation signal with the dark matter distribution having power law density profile $\rho(r) \propto r^{-\alpha}$, $\alpha \approx 1.34$. The observed excess is at energies between $\sim 600$~MeV and $\sim 6$~GeV and the mass of the proposed DM particle was suggested to be in the GeV energy band. In this work we analyze the \fermi\ data, used in \citet{Hooper10}, utilizing the data analysis tool, provided by the \fermi\ team.

We conclude that the signal within central $1^\circ{-}2^\circ$, containing the ``excess'' found by \citealt{Hooper10} (\textbf{HG10} hereafter), can be well described by our model : (point sources plus Galactic and extragalactic diffuse background components). The discrepancy is then due to a different interpretation of the data. The spectrum of the central point source (1FGL J1745.6-2900c, probably associated with the Galactic black hole Sgr A$^*$) was taken in HG10 to be a featureless power-law starting from energies about 10~TeV (results of HESS measurements, blue points with error bars in Fig.~\ref{fig:gc_spectrum}, \citep{Aharonian:04,vanEldik:07}) and continuing all the way down to $\sim 1$~GeV. The flux attributed in this way to the central point source is significantly weaker than in the previous works. For comparison, the (PSF corrected) spectrum of the GC point source reported in~\citet{Chernyakova10} is shown in Fig.~\ref{fig:gc_spectrum} in green points. Its spectral characteristics are fully consistent with the results of 11-months \fermi\ catalog~\cite{1FGLcat} ($\sim 6\times 10^{-8}~\mathrm{cts/cm^2/s}$ above 1~GeV, compared to the $\sim 5\times 10^{-9}~\mathrm{cts/cm^2/s}$ at the same energies in HG10). The change of the slope of the source spectrum below $\sim 100$~GeV, as compared with the HESS data is explained by \citet{Chernyakova10} with the model of energy dependent diffusion of protons in the few central parsecs around the GC. Alternatively, the spectrum can be explained with the model developed in \citet{Aharonian05}. The low-energy (GeV) component of the spectra in this model is explained by synchrotron emission from accelerated electrons, while high-energy (TeV) one by inverse Compton radiation of the same particles. According to the analysis of \citet{1FGLcat,Chernyakova10} the central point source provides significant contribution to the flux in the 1.25$^\circ$ central region. HG10 suggest, apparently, a different interpretation. They assume that there is no significant change in the spectrum of the central source at $\sim 100$~GeV and the spectrum observed by HESS at high energies continues to lower energies. Then, large fraction of the flux between the energies $\sim 600$ MeV and $\sim 6$~GeV has to be attributed to the ``DM excess''. One of the reasons in favor of such an interpretation could be the feature in the total spectrum from the central region (rise between $\sim 600$~MeV and several GeV) discussed in HG10. Such a feature would also be consistent with a possible contribution from millisecond pulsars \citep{Abazajian10a}, that is also expected to have a maximum at $\sim 2-3$~GeV. To illustrate the nature of the spectral shape at these energies we collected ``front converted'' (\textsc{front}) photons from the region of the width $5^\circ$ around the Galactic Plane (the ``\textit{inner}'' region) and from the ``\textit{outer}'' region as demonstrated on the left panel in Fig.~\ref{fig:aeff}. The count rate from each of these regions was divided by the \emph{constant} effective area ($3500~\mathrm{cm}^2$) to obtain the flux.\footnote{The effective area of \fermi\ LAT is strongly energy dependent. The number $3500~\mathrm{cm}^2$, roughly corresponding to the effective area at $\sim 1$~GeV, is used here as a quick expedient (see below).} One sees that the total emission from both regions demonstrates the same spectral behavior as the excess of HG10, suggesting that this spectral shape is \emph{not} related to the physics of the several central degrees. This drop of flux at low energies is mainly due to the decreasing effective area of the satellite.\footnote{\url{http://www-glast.slac.stanford.edu/software/IS/glast_lat_performance.htm}} If we properly take into account the dependence of the effective area on energy, we obtain the spectrum that ``flattens'' at small energies and exceeds by a significant factor the flux from the central point source (as it should) (compare red and magenta points on the right panel in Fig.~\ref{fig:aeff}). Another reason for the decrease of the HG10 spectrum is the increase of \fermi\ LAT PSF at low ($\lesssim 1$ GeV) energies.\footnote{For example, for normal incidence 95\% of the photons at $1$ GeV are contained within $\sim 1.6^\circ$ and in $2.8^\circ$ at $500$ MeV} This means that if one collects photons from a relatively small region, such that a contribution from its boundary (with the PSF width) is comparable to the flux from the whole region, the spectrum would artificially decline, due to increasing loss of photons at low energies. To disentangle properly what photons in the PSF region had originated from a localized source, and what are parts of the diffuse background, special modeling is needed. In the monotonic spectrum of the GC, obtained by~\citet{Chernyakova10} both these effects (effective area and PSF) were taken into account as it was obtained from $10^\circ\times10^\circ$ region, using the \fermi\ software. To further check the nature of the emission from the central several degrees, we took a fiducial model, that contained the same galactic and extragalactic diffuse components plus all the same point sources, but \emph{excluding the point source in the center}. We then fit our data to this new model. Such a fit attempts to attribute as many photons as possible from the region around the GC to the emission of diffuse components. The procedure leaves strong positive residuals within the central $1{-}2^\circ$. The spectrum of these residuals is consistent with the spectrum of the central point source of~\citet{Chernyakova10} (green points in Fig.~\ref{fig:gc_spectrum}). To demonstrate, that the spatial distribution of these residuals is fully consistent with the PSF of \fermi, we compare their radial distribution in various energy bins with the radial distribution around the Crab pulsar (as it was done e.g. in~\citet{Neronov:10b}). The pulsar wind nebula, associated with the Crab has an angular size $\sim 0.05^\circ$~\citep{Hester:08}. Thus, for \fermi\ LAT Crab is a point source. The radial profile of residuals at all energies has the same shape as Crab, as Fig.~\ref{fig:crab} clearly demonstrates. As an additional check, we repeated the above test using only \textsc{front} photons (as in this case the PSF is more narrow) and arrived to the same conclusion. The above analysis demonstrates that the emission around the GC in excess of diffuse components (galactic and extragalactic) is fully consistent with being produced by the point source with the power-law spectrum, obtained in~\cite{1FGLcat,Chernyakova10}, \emph{and no additional component is required.} \begin{figure*} \centering \begin{minipage}{.5\textwidth} \includegraphics[width=\linewidth]{1000MeV_front_profile} \end{minipage}~\begin{minipage}{.5\textwidth} \includegraphics[width=\linewidth]{1500MeV_front_profile} \end{minipage} \begin{minipage}{.5\textwidth} \includegraphics[width=\linewidth]{2500MeV_front_profile} \end{minipage}~\begin{minipage}{.5\textwidth} \includegraphics[width=\linewidth]{5000MeV_front_profile} \end{minipage} \caption{Radial profile of residuals at different energies around the GC as compared to the radial profile of Crab emission (renormalized so that the total flux in each energy range coincide). In both cases only \textsc{front} photons were used.} \label{fig:crab} \end{figure*} A different question however is whether such an additional component may be ruled out. To this end we have added to our model of Sec.\ref{sec:Model} an additional spherically symmetric component, whose intensity is distributed around the center as $\rho^2(r)$ (where $\rho(r) \propto r^{-1.34}$, as found in HG10). We observe, that such a procedure does improve the fit (change in the log-likelihood is 25 with only one new parameter added). The resulting spectral component is shown in Fig.~\ref{fig:dm}. Some of the photons from the galactic diffuse background were attributed by the fit procedure to the new component, concentrated in several central degrees (within the Galactic Plane). This phenomenon is probably related to the complicated and highly non-uniform in the central region galactic diffuse background\footnote{See ``\textit{Description and Caveats for the LAT Team Model of Diffuse Gamma-Ray Emission}'' by the Diffuse and Molecular Clouds Science Working Group, Fermi LAT Collaboration, \url{http://fermi.gsfc.nasa.gov/ssc/data/access/lat/ring_for_FSSC_final4.pdf}.} (cf. also the right panel of the Fig.~\ref{fig:model_cnt_map}). \begin{figure} \centering \includegraphics[width=\linewidth]{DM2_5_2} \caption{Spectrum of an additional spherically symmetric component, distributed around the GC as the HG10 excess.} \label{fig:dm} \end{figure} We should also note that HG10 modeled diffuse background differently. They considered contributions from the Galactic disk and spherically symmetric emission in the region \emph{outside} central $2^\circ$ and then extrapolated the diffuse model into the innermost $1^\circ-2^\circ$, arguing that the contribution does not vary significantly in the range $2^\circ - 10^\circ$ off-center. The background model we used (see \citealt{1FGLcat,FermiEGB} for the detailed description) is different from that of HG10, especially in the central 1-2$^\circ$, where the model flux is higher than the one extrapolated from larger galactic longitudes, as one can clearly see on the right panel of the Fig.~\ref{fig:model_cnt_map}. \begin{figure*} \includegraphics[width=.45\textwidth, angle=-90]{model_cntmap} \caption{Left: 10$^\circ$x10$^\circ$ count map of best-fit model. Right: only contribution from galactic and extragalactic backgrounds is shown} \label{fig:model_cnt_map} \end{figure*} \bigskip Having the above considerations in mind, we think that the spectrum of the central region, changing monotonously with the energy, is well described by purely astrophysical model of the central point source and therefore present data do not require any additional physical ingredients, such as DM annihilation signal or % additional contributions from millisecond pulsars. However, to firmly rule out the emission from DM annihilation in the GC, more detailed model of the galactic diffuse background is required. Additionally, with the future data, better statistics will reduce the error bars on the data point around $\sim 100$~GeV which will be helpful to better understand the central point source physics. \subsubsection*

1012.5839

1012

1012.0307_arXiv.txt

High energy emission from blazars is thought to arise in a relativistic jet launched by a supermassive black hole. The emission site must be far from the hole and the jet relativistic, in order to avoid absorption of the photons. In extreme cases, rapid variability of the emission suggests that structures of length-scale smaller than the gravitational radius of the central black hole are imprinted on the jet as it is launched, and modulate the radiation released after it has been accelerated to high Lorentz factor. We propose a mechanism which can account for the acceleration of the jet, and for the rapid variability of the radiation, based on the propagation characteristics of large-amplitude waves in charge-starved, polar jets. Using a two-fluid ($e^\pm$) description, we find the outflows exhibit a delayed acceleration phase, that starts when the inertia associated with the wave currents becomes important. The fluids propagate with the wave at approximately the sonic speed, corresponding to a bulk Lorentz factor $\gamma\approx10^4 \Delta t_{100}^{1/3}\kappa_{r_{\rm g}}^{-1/3}L_{46}^{1/6}M_9^{-1/3}$ out to radius $r_1\approx\Delta t_{100}^{1/3} \kappa_{r_{\rm g}}^{2/3}L_{46}^{1/6}M_9^{2/3}\, \textrm{pc}$, after which the Lorentz factor accelerates as $\gamma\propto r$. ($\Delta t_{100}$ is the variability time in units of $100\,$s, $\kappa_{r_{\rm g}}$ the pair multiplicity at one gravitational radius, $L_{46}$ the \lq\lq $4\pi$-luminosity\rq\rq\ of the jet in units of $10^{46}\,\textrm{erg/s}$, and $M_9$ the black-hole mass in units of $10^9\,\msolar$.) The time-structure imprinted on the jet at launch modulates photons produced by the accelerating jet provided $\kappa_{r_{\rm g}}<14\, \Delta t_{100} L_{46}^{1/8}M_{9}^{-1}$, suggesting that very rapid variability is confined to sources in which the electromagnetic cascade in the black-hole magnetosphere is not prolific.

Observations by the H.E.S.S. collaboration of TeV gamma-ray emission from the blazar PKS~2155-304 reveal very rapid variability \citep{HESS_2155a_07,HESS_2155b_10}, at very high flux levels. In the most extreme flare, variations on a timescale of a few hundred seconds at a flux level corresponding to an isotropic luminosity of $10^{46}\,\textrm{erg/s}$ were measured. If, as expected, the mass $M$ of the central black hole is $2\times10^9\,\msolar$, these observations imply structure in the jet that is roughly one hundred times smaller than the gravitational radius $r_{\rm g}=GM/c^2$ \citep{begelmanfabianrees08}. Though this is the most extreme example, several other blazars exhibit very rapid variability in GeV and TeV gamma-rays \citep[e.g.,][]{albertetal07,ackermannetal10}, which it is proving difficult to accommodate in the standard synchrotron-self-Compton picture, mainly because of the very high Lorentz factor and low magnetisation required of the jet \citep{levinson07,boutelierhenripetrucci08,graffetal08,katarzynskietal08,kusunosetakahara08,mastichiadismoraitis08,neronovsimikozsibiryakov08,ghisellinietal09,gianniosuzdenskybegelman09,paggietal09,tammiduffy09,riegervolpe10,nalewajkoetal10}. In the framework of ideal MHD, radial (uncollimated) relativistic jets do not accelerate after they pass through the fast magnetosonic point. Collimation, however, requires special boundary conditions \citep{lyubarsky09,lyubarsky10}. It is, therefore, difficult to envisage the production of a jet with high Lorentz factor and low magnetisation. An isolated, impulsive, ejection event in an ideal MHD flow is able to circumvent this problem \citep{granotkomissarovspitkovsky10}, but, in a fluctuating jet, dispersion filters out the small timescale structure, and limits the acceleration \citep{levinson10}. It appears, therefore, that it may be necessary to go beyond the ideal MHD approximation in order to understand the observations of very rapid variability in blazars. Non-ideal MHD effects can become important when plasma is {\em charge-starved}: If the number density of charged particles is limited, this places a maximum on the absolute value of the charge density, which can, for example, result in the inability of the plasma to screen out the component of the electric field parallel to the magnetic field. Charge starvation also places an upper limit on the available current density. As this limit is approached, the relative drift-speed of the charged components becomes relativistic, and their associated inertia begins to contribute to the fluid stress-energy tensor. This situation might plausibly arise in a black-hole magnetosphere, because axisymmetric, general-relativistic, MHD simulations of jets launched by accreting, rotating black holes \citep{devilliershawley03,mckinneygammie04} reveal a conical region around the rotation axis into which the accreting plasma does not penetrate. As in a pulsar magnetosphere, the matter density in this region is likely to be determined not by accretion, but by the rate at which electron-positron pairs are created in the strong electromagnetic fields that penetrate it \citep{goldreichjulian69,blandfordznajek77,levinson00}. If an outflow results, the density of pairs decreases, and, far from the hole, non-MHD effects connected with particle inertia can become important. In the case of an axisymmetric, force-free magnetosphere, it is known that plasma is ejected, carrying off energy mainly in the form of Poynting flux via the \lq\lq Blandford-Znajek\rq\rq\ mechanism. On the axis itself, the energy flux vanishes, so that the polar regions of the jet do not dominate the overall energetics. However, observations of rapid variability suggest that axisymmetry may not be a good approximation, since they imply small-scale structure in the black-hole magnetosphere. Simulations in which the black-hole spin and the asymptotic magnetic field are misaligned find a Poynting flux comparable to the Blandford-Znajek value \citep{palenzuelaetal10}, making it plausible that also more complex non-axisymmetric field structures can power a substantial, magnetically dominated, polar jet. In the following, we develop this idea by examining the propagation characteristics of nonlinear electromagnetic waves above the polar regions of a rotating black hole. Using a model consisting of cold electron and positron fluids, we show that a circularly polarised magnetic shear propagates radially outwards at roughly the fast magnetosonic speed until it reaches the point where the ideal MHD description loses its validity. In (non-ideal)~MHD language, this happens when the inertia of the plasma particles begins to affect the conductivity, i.e., when the drift speeds implied by the plasma current become relativistic \citep{melatosmelrose96,meier04}. This can occur at a large distance from the black hole, depending on the density of injected pairs. The wave then goes through a phase of delayed acceleration in which it converts Poynting flux into kinetic energy flux. If particles radiate gamma-rays in the acceleration zone, then, despite the large spatial extent of the source, the spatio-temporal structure of the shear wave, that is imprinted on it close to the black hole, modulates the radiation, provided the mass-loading of the jet is sufficiently small. In section~\ref{parameters} we discuss the parameters used to specify the physical conditions in the jet. The two-fluid jet model is presented in section~\ref{two-fluid}. First, the nonlinear plane-wave solution representing a magnetic shear is derived, then radial propagation in spherical geometry is discussed. It is shown in the Appendix that general relativistic effects drop out in the short-wavelength approximation ($c/\omega\ll r$) when the Kerr metric is used. The application to blazar variability is discussed in section~\ref{blazars}, and our conclusions summarised in \ref{conclusions}.

\label{conclusions} In this paper, we describe a mechanism that causes a magnetically dominated, radial outflow from a black-hole magnetosphere to enter a delayed acceleration phase, starting at a distance from the hole given by (\ref{defracc}). Applying this mechanism to blazar jets, we derive a constraint, (\ref{limit}), on the pair density in the magnetosphere that would allow radiation produced where the jet accelerates to retain any short-timescale structure imposed on it close to the launching site. The mechanism is based on an analysis of the propagation characteristics of a nonlinear wave -- specifically a circularly polarised magnetic shear -- in a low-density plasma. Such a wave, we suggest, is likely to be launched in the polar regions of a rotating, accreting black hole, and, in a non-axisymmetric picture, may fluctuate on a time shorter than $r_{\rm g}/c$, as indicated by observations of the source PKS~2155-304. Acceleration is a result of charge-starvation -- a non-MHD effect that arises when the relative drift-speed of the oppositely charged constituents in a low-density plasma becomes relativistic. The analysis employs a cold two-fluid model of the plasma, and uses a short-wavelength perturbation expansion to find the evolution of the radially propagating, nonlinear wave. The equations are derived in Kerr geometry. However, under the conditions we envisage, where the wavelength of the oscillation is of the same order in the expansion parameter as the gravitational radius, general relativistic effects do not appear in the governing equations. Several important problems remain to be investigated. These include the nature of the dissipation and radiation mechanisms, and the effect these might have on the propagation of the wave, as well as the possibility of modelling the multi-wavelength blazar spectrum. Furthermore, although the picture of a circularly polarised magnetic shear that is static in the jet frame is intuitively attractive, this is only one specific, nonlinear solution of the governing equations; other polarisations and other modes, such as the linearly polarised \lq\lq striped wind\rq\rq\ \citep{lyubarskykirk01} or the electromagnetic mode of superluminal phase-speed \citep{kirk10} may also prove important. Nevertheless, the underlying physical cause of the acceleration --- the inertia of the charge-carriers --- suggests that delayed jet-acceleration may be a generic phenomenon. \appendix

1012.0307

1012

1012.1137_arXiv.txt

The CO$_2$ production rate has been derived in comets using the Cameron band (a$^3\Pi$ - X$^1\Sigma$) emission of CO molecule assuming that photodissociative excitation of CO$_2$ is the main production mechanism of CO in a$^3\Pi$ metastable state. We have devoloped a model for the production and loss of CO(a$^3\Pi$) which has been applied to comet 103P/Hartley 2: the target of EPOXI mission. Our model calculations show that photoelectron impact excitation of CO and dissociative excitation of CO$_2$ can together contribute about 60-90$\%$ to the Cameron band emission. The modeled brightness of (0-0) Cameron band emission on comet Hartley 2 is consistent with Hubble Space Telescope observations for 3-5\% CO$_2$ (depending on model input solar flux) and 0.5\% CO relative to water, where photoelectron impact contribution is about 50-75\%. We suggest that estimation of CO$_2$ abundances on comets using Cameron band emission may be reconsidered. We predict the height integrated column brightness of Cameron band of $~\sim$1300 R during EPOXI mission encounter period.

\label{sect1} In the exploration of the solar system, comets have been targeted by various space missions. After successful encounter of comet 9P/Temple 1 on 4 July 2005, the NASA's Deep impact mission, also called EPOXI mission, under it's extended investigation program DIXI will encounter comet 103P/Hartley 2 on 4 November 2010, with closest approach around 700 km from the nucleus. This comet has been observed by several space telescopes in different spectral regions \citep{Weaver94,Crovisier99,Colangeli99,Groussin04,Lisse09,Snodgrass08,Snodgrass10}. The first clear detection of the Cameron band (a$^3\Pi$ - X$^1\Sigma$) of CO was reported by \cite{Weaver94} in HST/FOS spectrum of comet 103 P/Hartley 2. Since Cameron band emission is a forbidden transition, resonance fluorescence is not an effective excitation mechanism. The upper state of this emission (a$^3\Pi$) is a metastable state with lifetime of 3 ms \citep{Gilijamse07}, which is quite small. Thus, Cameron band emissions are treated as ``prompt emissions'' and can be used to probe distribution of parent species if this emission is produced in dissociative excitation of a molecular species. Photodissociative excitation of CO$_2$ is considered as the major production mechanism of CO Cameron band and has been used to trace the distribution and abundance of CO$_2$ on comets \citep{Feldman97,Weaver94}. Besides photodissociative excitation of CO$_2$, there are other channels of excitation of CO molecule in the a$^3\Pi$ state. It has been shown that photoelectrons generated by solar EUV radiation also play an important role in excitation, dissociation, and ionization processes leading to emission and chemistry in cometary comae \citep[e.g.,][]{Ip86,Boice86,Korosmezey87,Bhardwaj90,Bhardwaj96,Bhardwaj99a,Bhardwaj03,Weaver94, Haider05,Capria08}. Recently, \cite{Campbell09} demonstrated the importance of photoelectron impact excitation in comets, and showed that electron impact on CO gives 40\% contribution to the total CO Fourth positive emission. The presence of photoelectron excitation in cometary coma is clearly publicised by detection of OI 1356 \AA \ emission in comets \citep[e.g.,][]{Sahnow93,Mcphate99}, since this emission being a spin-forbidden transition cannot be produced by solar fluorescence. There are other siginificant evidences for an important role of photoelectron excitation in cometary coma \citep[e.g.,][]{Tozzi98,Bhardwaj99a,Feldman09}. In addition to photon and electron impact reactions, dissociative electron recombination reactions of CO$^+$-bearing ions can also produce CO in the a$^3\Pi$ excited state. Our aim in this paper is to study various production and loss mechanisms of CO(a$^3\Pi$) and to estimate the contribution of photoelectron impact excitation of CO and CO$_2$ in the production of Cameron band for different relative abundances of CO$_2$ on comet 103 P/Hartley 2: the target of EPOXI mission. Since model calculations depend on input solar flux, we have estimated its sensitivity on the calculated intensity of Cameron band emission. We show that photoelectron impact on CO and CO$_2$ are dominant processes ($\sim$60--90\% contribution) in producing CO molecule in (a$^3\Pi$) state. Around the EPOXI encounter epoch predictions are made for the brightness of Cameron band on for comet 103P/Hartley 2, which will be observed by several space-based telescopes including HST. \section {Model} \label{sect2} We have developed a model for the production of Cameron band emission on comets, which uses the basic coupled chemistry model described in detail in our earlier papers \citep{Bhardwaj96, Bhardwaj99a, Haider05}. Various sources of production and loss of CO(a$^3\Pi$) are summarized in Table \ref{prodlosstab}. The total water production rate is taken as 6.3 $\times$ 10$^{28}$ s$^{-1}$ for comet 103P/Hartley 2 \citep{Weaver94}. To evaluate the effect of solar EUV flux on model calculations, we have considered 2 solar flux models: EUVAC model of \cite{Richards94} and SOLAR 2000 (S2K) model of \cite{Tobiska00}. The degradation of the solar UV-EUV radiation and solar EUV-generated photoelectrons in the coma is modeled using the method of \cite{Bhardwaj90,Bhardwaj96} and further developed by \cite{Bhardwaj99a,Bhardwaj03}. The electron impact production rates are calculated using the Analytical Yield Spectrum (AYS) approach, which is based on the Monte Carlo method. Details of AYS approach are given in several of the previous papers \citep{Bhardwaj90,Bhardwaj96, Bhardwaj99a,Bhardwaj99d, Bhardwaj99b,Bhardwaj09}. The current model takes into account the most recently published cross sections for the photon impact and electron impact dissociation, ionization, and excitation processes for the gases in the coma. The cross section for photodissociative excitation of CO$_2$ producing CO in a$^3\Pi$ state is calculated using absorption cross sections of CO$_2$ and yield of Cameron band measured by \cite{Lawrence72}. The cross section for electron impact excitation of CO(a$^3\Pi$) from CO is taken from \cite{Jackman77} and for dissociative excitation of CO$_2$ is taken from \cite{Bhardwaj09}. The electron temperature profile required for dissociative recombination reactions is taken from \cite{Korosmezey87} and is assumed to be same as on comet Halley. Calculations are made for the comet 103P/Hartley 2 at heliocentric distance of 0.96 AU.

\label{sect3} The photodissociation of CO$_2$ producing CO in a$^3\Pi$ state is determined by solar flux mainly in the wavelength region 550 to 1050 \AA. Table~\ref{a3piprodtab} presents the calculated photon production frequencies of CO(a$^3\Pi$) for two different solar flux models. The CO(a$^3\Pi$) production frequencies calculated for photoelectron impact on CO$_2$ and CO are also shown in the same table for the corresponding solar flux models. Our calculated photodissociation frequencies are about 50\% to a factor of 2 lower than those reported by \cite{Huebner92}. Using EUVAC solar flux, the calculated radial profile of volume production rate for the various sources of CO(a$^3\Pi$) at the relative abundance of 4\% CO$_2$ and 0.5\% CO are shown in Figure~\ref{proda3pi}. At 100 km cometocentric distance, the dominant source of production of CO(a$^3\Pi$) is electron impact of CO$_2$ ($\sim$50\%) followed by electron impact of CO ($\sim$25\%), and photodissociation of CO$_2$ ($\sim$15\%). The contributions from dissociative recombination reactions are quite small ($\le$5\%) at lower cometocentric distances, but the recombination of CO$_2^+$ is a significant ($<$30\%) source at 1000 km and beyond. Figure~\ref{lossa3pi} shows radial profile of various loss processes of CO(a$^3\Pi$) for the same relative composition of CO$_2$ and CO. Since lifetime of CO in excited state (a$^3\Pi$) is very short \citep[$\sim$3 ms;][]{Gilijamse07}, the radiative decay is the dominant loss process. Collisional quenching of CO(a$^3\Pi$) by cometary neutral species is negligible since 103P/Hartley 2 is a low production rate comet. But in the case of large production rate comets, like Hale-Bopp, quenching by water would be a dominant loss process in the innermost part of the coma. Figure~\ref{proja3pi} shows the modeled limb brightness profiles of Cameron band emission for different production processes of CO(a$^3\Pi$). The cometary coma is assumed to be spherically symmetric. The production rates are integrated up to 10$^5$ km along the line of sight at a given projected distances from the cometary nucleus, and converted into brightness. The brightness profiles are averaged over the projected area of slit (2870 $\times$ 954 km) corresponding to the HST observation \citep{Weaver94}. The volume emission rate for 3 transitions (0-0, 1-0, 0-1) of Cameron band emission are calculated using the following formula \begin{equation} V_{\nu'\nu''}(r)=q_{o\nu'} (A_{\nu'\nu''}/\sum_{\nu''}A_{\nu'\nu''})\ V(r)\ exp(-\tau) \end{equation} where V(r) is total volume excitation rate of CO(a$^3\Pi$) at cometocentric distance r, q$_{o\nu'}$ is the Franck-Condon factor for transition, A$_{\nu'\nu''}$ is Einstein transition probability from upper state $\nu'$ to lower state $\nu''$, and $\tau$ is optical depth. Since resonance fluorescence is not an effective excitation mechanism for the Cameron band and the total gas production rate is only 6.3 $\times$ 10$^{28}$ s$^{-1}$, the cometary coma can be safely assumed to be optically thin. The Franck-Condon factors are taken from \cite{Nicholls62} and branching ratios from \cite{Conway81}. The relative contributions of (1-0), (0-0), (0-1) transitions to the total Cameron band are 13.9\%, 10.4\%, and 14.7\%, respectively. Table~\ref{bigtab} presents the model calculated slit-averaged brightness of (1-0), (0-0), (0-1) transitions of Cameron band, as well as total Cameron band brightness and height-integrated column brightness for different relative abundances of CO and CO$_2$ corresponding to the HST observation of comet 103P/Hartley 2 on September 18-19, 1991. Due to the absence of CO Fourth positive emission in this comet \citep{Weaver94}, the abundance of CO is constrained to 0.5\%. However, we do consider a case of 1\% of CO to evaluate it's implications on the results. This table also depicts fractional contribution of photodissociation of CO$_2$, photoelectron impact of CO and CO$_2$, and dissociative recombination of CO$_2^+$ to the total calculated brightness at 3 projected distances (10$^2$, 10$^3$, and 10$^4$ km) from the nucleus. Since the production rates of photodissociative excitation of CO$_2$, and photoelectron impact of CO and CO$_2$, are dependent on input solar flux model, results are presented for the EUVAC and S2K solar fluxes relevant to the date of comet observation which was in solar maximum condition. The HST observation of 0-0 transition of Cameron band is 35 Rayleigh \citep{Weaver94}, which is consistent with model calculated brightness for the relative abundance of 4 to 5\% of CO$_2$ and 0.5\% CO when EUVAC solar flux is used. In this case, at 100 km, the photoelectron impact of CO$_2$ (50\%) and CO (25\%) contribute around 75\%, while photodissociative excitation of CO$_2$ is $<$15\%. At 1000 km and beyond, the contribution due to electron impact of CO$_2$ and CO is about 60--70\% while those of dissociative recombination of CO$_2^+$ is $\sim$15--30\% and of photodissociative excitation of CO$_2$ $\sim$10\% only. On an average, the photoelectron impact of CO$_2$ and CO contribute about 60-75\% to the production of Cameron band emission, while photodissociative excitation of CO$_2$ contribute about 10-15\% only. In the case of S2K solar flux model, the CO$_2$ abundance of 3 to 4\% is required to match HST-observed Cameron band 0-0 transition brightness. Here the contribution of photodissociative excitation of CO$_2$ is $\sim$20\%, while the electron impact of CO$_2$ and CO together contribute $\sim$65\%, to the total Cameron band emission. When the CO abundance is doubled to 1\% of water the relative contribution due to electron impact on CO increases resulting in the reduction (by $\sim$1\%) in the requirement for CO$_2$ abundace to match the HST-observation brightness. However, there is no major change in the percentage contributions due to photodissociation and photoelectron impact excitation of CO and CO$_2$. Hence, we conclude that the photodissociative excitation of CO$_2$ is not the dominant source for the production of Cameron band in comets.

1012.1137

1012

1012.3418_arXiv.txt

{ The ability to predict the future behavior of solar activity has become of extreme importance due to its effect on the near Earth environment. Predictions of both the amplitude and timing of the next solar cycle will assist in estimating the various consequences of Space Weather. The level of solar activity is usually expressed by international sunspot number ($R_z$). Several prediction techniques have been applied and have achieved varying degrees of success in the domain of solar activity prediction. In this paper, we predict a solar index ($R_z$) in solar cycle 24 by using the neural network method. The neural network technique is used to analyze the time series of solar activity. According to our predictions of yearly sunspot number, the maximum of cycle 24 will occur in the year 2013 and will have an annual mean sunspot number of 65. Finally, we discuss our results in order to compare it with other suggested predictions.

% \label{sect:intro} The successful prediction of a future event is arguably the most powerful way of confirming a scientific theory. Commonly in physics, a theory that is describing a system in a natural world is regarded as correct and therefore useful if it can use the state of the system at one time to reconstruct the state of the system at some other time, in past or future. The prediction of solar activity for a few years is the oldest problem in solar physics, arising as soon as solar cycle itself was discovered. Unfortunately, this problem has not been solved, probably because the series of observational data available are not long enough for purely statistical analysis, and because we do not quite understand the physical nature of this phenomena. Most of the space weather phenomena are influenced by variations in solar activity. During the years of solar maximum there are more solar flares causing significant increase in solar cosmic ray intensity. The high-energy particles disturb communication systems and affect the lifetime of satellites. Coronal mass ejections and solar flares are the origin of shocks in solar wind and cause geomagnetic disturbances in the earth's magnetosphere. The high rate of geomagnetic storms and sub-storms results in atmosphere heating and drag of Low Earth Orbit (LEO) satellites. Solar activity forecasting is especially useful to space mission centers as in the orbital trajectory parameters of satellites are greatly affected by variations of solar activity. A dramatic effect, not only on the Earth's upper atmosphere, disturbing the orbits of satellites, but also on power grids on the ground, e.g. the power cuts in Quebec, Canada in 1989. The level of solar activity is usually expressed by the Zurich or International sunspot number. Although the solar activity presents some clear periodicities, its prediction is quiet difficult but not impossible, as a large range of forecasting methods using predict the occurrence and amplitude of solar cycle is categorized to two models; statistical models and physical models. In statistical models, it is usual to represent the evolution of a physical system by using a time series. In Contrast with a physical model, the statistical model only attempts to explain the system, and in particular a time series associated with it, in terms of itself, and perhaps in terms of correlation with other time series associated with the system. At this point, it is appropriate to address a common concern, which for obvious reasons is most usually expressed by physicists: what reasons are there for constructing a model that contains no physical understanding? Here are three reasons. Firstly, simply writing down the data as a time series, together with organizing and examining it, is the first step in the scientific method: analyzing the sequence as a time series governed by a statistical model is a natural first step, until such time as a physical theory can be formulated. The second reason is that predictions from a statistical model might simply be useful in their own right. For example, in day to day life it makes no difference to most people whether the weather forecast was made from a statistical model or from a physical one. The final and most important reason is that it might be impossible for the physical system to be predicted from the basic physical principles governing it. This can be because the system is simply too complicated, which, for example, is the case for a plasma (Conway~\cite{con98}). One of the statistical models using for predicting the data is artificial neural networks method. The use of artificial neural networks has been recognized recently as a promising way of making predictions on temporal series with chaotic or irregular behavior (Weigend~\cite{wei90}). This technique has already been applied in the framework of solar-terrestrial physics for prediction of geomagnetic induced current and storms (Lundstedt~\cite{lun92}) and as a way of recognizing a pattern in the onset of a new sunspot cycle (Koons~\cite{koo90}). The aim of this paper is to predict the solar cycle. The structure of the paper is as follow. In section 2 we provide a brief summary of the neural network methodology employed. In section 3 we introduce the results of our network architecture to generate our best estimate of the behavior of cycle 24, and In section 4, the conclusions and their comparison are presented.

\label{sect:discussion} Our neural network method is based on one hidden layer. For having reliable result, we use multi-step prediction to have only one reasonable output. In terms of processing data, by changing back-propagation algorithm to Levenberg-Marquardt algorithm our feed-forward neural network model becomes faster since LM algorithm speeds up convergence while limiting memory requirements (Battiti~\cite{bat92}). We saw almost a similarity between predicted Solar cycle 24 and Solar cycle 20 . We predict a SC 24 with a maximum of $65\pm13$ occurring in 2013. In general, our result is close to other prediction made for solar cycle 24. For example, Li et al.~\cite{li05} obtained 2013 a maximum of cycle 24 with statistical method. Also a recent article by Wang et al.~\cite{wang09}, using similar descending phases and a cycle grouping, predicted that peak amplitude for that monthly smoothed sunspot number in the solar cycle 24 is near $100.2\pm7.5$, occurring in 2012. Furthermore, Chumak et al.~\cite{chumak10} predicted that the maximum amplitude of cycle 24 is $90\pm20$ which is in agreement with our results. Finally, Our prediction fits well within the limits of the others as indicated in Pesnell~\cite{pes08} where an average cycle was predicted using other methods such as statistical and precursor methods.

1012.3418

1012

1012.1247_arXiv.txt

We examine the constraints on the luminosity-dependent density evolution model for the evolution of blazars given the observed spectrum of the diffuse gamma-ray background (DGRB), blazar source-count distribution, and the blazar spectral energy distribution sequence model, which relates the observed the blazar spectrum to its luminosity. We show that the DGRB observed by the Large Area Telescope (LAT) aboard the {\it Fermi Gamma Ray Space Telescope} can be produced entirely by gamma-ray emission from blazars and nonblazar active galactic nuclei, and that our blazar evolution model is consistent with and constrained by the spectrum of the DGRB and flux source-count distribution function of blazars observed by Fermi-LAT. Our results are consistent with previous work that used EGRET spectral data to forecast the Fermi-LAT DGRB. The model includes only three free parameters, and forecasts that $\gtrsim 95\%$ of the flux from blazars will be resolved into point sources by Fermi-LAT with 5 years of observation, with a corresponding reduction of the flux in the DGRB by a factor of $\sim$2 to 3 (95\% confidence level), which has implications for the Fermi-LAT's sensitivity to dark matter annihilation photons.

The source of the extragalactic isotropic diffuse gamma-ray background (DGRB) has been an unsolved question in astrophysics for some time. In this paper, we show how the DGRB spectrum can be produced by a combination of blazar and nonblazar active galactic nuclei (AGN) gamma-ray sources. We also show that the blazar flux source-count distribution function ($dN/dF$) is consistent with the full DGRB originating from these sources. Furthermore, we show how less-detailed models of the blazar contribution failed to be consistent with the DGRB. We explore how the implications for dark matter detection or constraints from the DGRB will evolve as the blazar sources of the DGRB are resolved. The DGRB was first discovered by the SAS 2 experiment in 1975, for gamma-ray emission in the range of 35 to 300 MeV~\cite{Fichtel75,*Fichtel77,*Fichtel78,*Thompson82}. This background was seen at energies up to 20 GeV by the EGRET Collaboration, and it was confirmed at these energies in the first-year data from the Large Area Telescope (LAT) aboard the {\it Fermi Gamma-Ray Space Telescope}~\cite{Sreekumar98,Abdo10b,Abdo10}. The assumed extragalactic source of the DGRB is determined by measuring the complete diffuse (unresolved) flux and then subtracting off a model to account for the background coming from our Galaxy. This yields a measure of the flux coming from unresolved diffuse sources, presuming there is no minimal isotropic component from the Galaxy, e.g. dark matter annihilation or decay. The DGRB has been used to constrain dark matter annihilation in Galactic and extragalactic sources~\cite{Abdo10c,Abazajian10,Baltz08}. The most recent measurement of the DGRB was performed by the Fermi-LAT. In the Fermi-LAT Collaboration analysis, the gamma-ray intensity was measured in the range of 100 MeV to 100 GeV above $10\degree$ in Galactic latitude ($\lvert b\rvert >10\degree$). The total flux is modeled by stacking the spectra of known sources with the cosmic-ray background, the Galactic diffuse background, and the DGRB. This analysis gives a DGRB intensity that is roughly 25\% of the total observed flux. The DGRB seen by the Fermi-LAT is consistent with a power law in energy with index 2.41. This value for the DGRB is notably softer at high energies than was previously seen in the EGRET Collaboration, which is partly due to an updated model of the diffuse Galactic emission in Ref.~\cite{Abdo10b} (hereafter FS10). A detailed spectral energy distribution (SED) sequence model of blazars can reproduce the DGRB~\cite{Inoue09,Inoue10a}. We explore this model in this work. Many models have been proposed to explain the DGRB. It has been shown that emission from AGN can account for the diffuse background from 10 keV to 100 MeV, but above that energy, this model cannot account for the large gamma-ray flux~\cite{Inoue08}. Radiation from star-forming galaxies could account for much of the DGRB up to 10 GeV, but this also cannot explain the high intensities observed at higher energies~\cite{Fields10}. Emission from millisecond pulsars has been proposed as a source as well~\cite{FaucherGiguere09}. However, millisecond pulsars as a dominant source of the DGRB may be inconsistent with the lack of anisotropy in the DGRB~\cite{SiegalGaskins10}. Dark matter annihilation, both as a component of the extragalactic diffuse emission and as an unaccounted foreground from the Milky Way can contribute to the DGRB, but the fluxes from dark matter are expected to be lower than the DGRB flux and have a different spectral shape~\cite{Abazajian10,Abdo10c}. However, measurements of the DGRB are one of the strongest ways to constrain dark matter annihilation~\cite{Baltz08}. If dark matter is a significant contributor, it may be disentangled from astrophysical sources due to its angular correlation on the sky~\cite{Ando:2005xg,*Ando:2006cr,*Miniati07,*SiegalGaskins:2008ge,*SiegalGaskins:2009ux,*Hensley09,*Fornasa:2009qh}. Pioneering work proposed that blazars could account for {\it all} of the DGRB seen by the EGRET Collaboration~\cite{Stecker96}. The blazar class of AGN has been studied in depth as the origin of the DGRB at high energies~\cite{Stecker93,Stecker96,Padovani93,Salamon94,Chiang95,Chiang98,Mucke00,Giommi05,Narumoto06,Dermer07,Pavlidou08,Bhattacharya09}. In Ref.~\cite{Inoue09} it was shown that the DGRB can be composed of blazars and nonblazar AGN in the luminosity-dependent density evolution (LDDE) SED blazar model. This model contains only three free parameters describing the gamma-ray luminosity function (GLF) of blazars. We show that this model is consistent with producing the full DGRB spectrum as well as the blazar source-count distribution, $dN/dF$, of blazars as measured by Fermi-LAT. In addition, we constrain this model by these measurements and find parameters for which the model successfully reproduces these measurements. Note that both the source-count distribution $dN/dF$ and DGRB spectrum are predicted by the model, and not an input to the model. Recent work by the Fermi-LAT Collaboration found that the DGRB could not be composed entirely by blazars~\cite{Abdo10} (hereafter FB10). However, that work adopted an over-simplification of the blazar SED to be a single power-law (PL), independent of blazar luminosity, which is inconsistent with the observed spectral luminosity dependence seen in the SED sequence~\cite{Fossati97,*Fossati98,*Donato01}. In contrast, in a separate paper, the Fermi-LAT Collaboration emphasizes the need for including departures from pure-PL behavior in blazar spectra when calculating the contribution of unresolved low-luminosity blazars to the DGRB~\cite{Abdo10a}. Incorporating the SED departure and its dependence on blazar luminosity evolution when modeling the DGRB is exactly the intent of the work presented here. Furthermore, the blazar model in FB10 lacks a physical evolution model for blazars. Instead of the source-count distribution resulting from the cosmological evolution of blazars, the source-count distribution is an input to the model, as a broken power-law with four free parameters. Note that even though the model in FB10 is simplistic, it contains {\it more} free parameters than the LDDE plus SED-sequence model explored here. In our approach there are three parameters in the adopted blazar model which describe the relation between the GLF and x-ray luminosity function (XLF). Because the FB10 model employs a pure-PL luminosity-independent SED with a broken-PL source-count distribution, the conclusions of that work do not apply to the model examined here. Other parameters in our work ({\it e.g.}, the SED sequence and the low-energy nonblazar AGN model) are constrained by other observations and remain fixed in our blazar model analysis. Namely, the observational constraints on the SED sequence come from spectral population models of blazars as in \cite{Fossati97,*Fossati98,*Donato01}, and the nonblazar AGN spectrum is constrained by the hard x-ray luminosity function derived from {\em HEAO1}, {\em ASCA}, and {\em Chandra} x-ray AGN surveys \cite{Ueda03,Inoue08}. A recent paper by Malyshev and Hogg~\cite{Malyshev:2011zi} using the one-point probability distribution function (PDF) of the DGRB also concludes that blazars cannot constitute the total DGRB flux as measured by Fermi-LAT, when modeled as a pure-PL SED with a fixed $dN/dF$. However, this conclusion also only applies to the model which they consider, which adopt blazars as having pure-PL luminosity-independent SEDs, and not to the LDDE SED-sequence model examined here. Because observed blazars make up about 15\% of the total gamma-ray flux, unresolved blazars are a likely candidate to make up the DGRB~\cite{Sreekumar98,Abdo10}. Blazars were the most numerous point-source objects observed by the EGRET Collaboration~\cite{Hartman99}. Additionally, observed blazar spectra tend to follow a similar power law in energy as the DGRB. However, it is known that blazars have a luminosity dependence to their spectral shape, which is incorporated in the SED-sequence model \cite{Fossati97,*Fossati98,*Donato01}, but ignored in the analysis of FB10. Blazars are the combination of two classes of AGN: flat-spectrum radio quasars (FSRQs) and BL Lacertae objects (BL Lacs). FSRQs are AGN that have spectral index $\alpha_{r}<0.5$ in the radio band and have radio emission lines with equivalent width greater than $5\rm\ \AA$. BL Lacs have no strong absorption or emission features, and have equivalent widths less than $5\rm\ \AA$~\cite{Urry95}. Broadly speaking, blazars tend to have their bolometric luminosities dominated by the gamma-ray luminosity and have great variability in that luminosity. Therefore, it is believed that blazars represent the small set of AGN that are observed along the jet axis, as opposed to nonblazar AGN which are observed far from the jet axis and dominate emission by their luminous accretion disk. This jet source is expected to be relativistically beamed, as opposed to the more isotropic flux coming from the AGN's accretion disk~\cite{Blandford79,Dermer95}. Different models of blazar emission have been proposed in the literature~\cite{Padovani93,Stecker93,Salamon94,Stecker96,Mucke00,Giommi05,Pavlidou08,Chiang95,Chiang98,Dermer07,Bhattacharya09,Narumoto06}. One is the pure luminosity evolution (PLE) model of the distribution of blazars~\cite{Chiang95,Chiang98,Dermer07,Bhattacharya09}. In this model, only the blazar luminosity is evolved in redshift. An alternative model, LDDE, relates the gamma-ray luminosity of blazars to the redshift-dependent distribution of x-ray emission from nonblazar AGN~\cite{Narumoto06}. This technique more realistically fits the blazar evolution to the AGN distribution, rather than assuming that all blazars have identical evolution regardless of luminosity. In many models for blazar spectra, a simple power-law or distribution of power laws is used as the intrinsic blazar spectrum, but more detailed frequency-dependent models have been used as well~\cite{Giommi05}. Here, we employ the LDDE model for blazar distributions. For the intrinsic spectrum of blazars, we use a frequency-dependent SED based on the multiwavelength study of Ref.~\cite{Fossati97,*Fossati98,*Donato01}. We use these models to derive the differential blazar spectrum in redshift, luminosity, and energy. By integrating over these variables, we can determine the number of detectable blazars for given detector sensitivities, and we can calculate the expected gamma-ray flux from unobserved blazars to determine how significantly they contribute to the DGRB. Additionally, we add a fixed nonblazar AGN component to our predicted blazar flux, which should make the net flux from our model fit the diffuse background over the energy range from 10 keV to 100 GeV. Below, we begin by describing the DGRB seen by the Fermi-LAT as well as its data on blazars. We will then describe our model in detail, specifying the evolution model and SED used in our calculations and how we fit these to the known data. We use this model to predict the ability of the Fermi-LAT to detect blazars and how this will affect the DGRB. Throughout the paper, we take a flat universe with the cosmological parameters $\Omega_{m}=0.272$, $\Omega_{\Lambda}=0.728$, and $H_0=70.2\rm\ km\ s^{-1}\ Mpc^{-1}$~\cite{Komatsu10}. {\it Note,} the use of $h$ in the text refers to Planck's constant, and not the Hubble parameter.

We have shown that the DGRB can be composed entirely by gamma-rays produced in blazars and nonblazar AGN. The LDDE plus SED-sequence is a physical model for the spectral evolution of a cosmologically-evolving blazar population contributing to the DGRB based on the unified AGN model for blazars. This model successfully accounts for the full DGRB spectrum as well as the full blazar source-count distribution function, which, unlike other approaches, are not used as components of the model. Independent of the nonblazar AGN component, the blazar model produces nearly the entire DGRB at its highest measured energies. The small value of $\kappa\simeq 2.4\times10^{-6}$, the x-ray AGN fraction seen as blazars, constrains this model to require a small fraction, $\lesssim$20\%, to be both properly oriented and sufficiently energetic in order to be gamma-ray emitters. We found constraints on this model from the spectrum of the DGRB and source-count distribution function $dN/dF$ of blazars as observed by Fermi-LAT. Our results are consistent with previous work by Inoue \& Totani \cite{Inoue09} which employed EGRET spectral data to forecast the Fermi-LAT DGRB. We forecast that $94.7^{+1.9}_{-2.1}\%$ of the flux from blazars will be resolved into point sources by Fermi-LAT with 5 years of observation, with a corresponding reduction of the flux in the DGRB by a factor of $\sim$2 to 3 (95\% CL) from the automatic removal of these sources in the measurement of the DGRB. This has significant consequences for the sensitivity of the DGRB measurement to dark matter annihilation, which we explore in a companion paper~\cite{Abazajian10a}. We predict that $2415^{+240}_{-420}$ blazars should be resolved, of $5.4^{+1.8}_{-1.7}\times 10^4$ total blazars in the universe (95\% CL). Recent results of anisotropy in the DGRB also indicate the likely presence of an unresolved point-source population~\cite{Vargas:2010en}. Using tests with enhanced point-source sensitivity, we find that future gamma-ray experiments at Fermi-LAT energies will resolve the blazar contribution to the DGRB such that the flux in the DGRB decreases as the square root of the point-source sensitivity. The LDDE plus SED-sequence model is more complex than the over-simplistic source-count method with a fixed spectral-index distribution adopted by the Fermi-LAT Collaboration in FB10, yet it has {\it fewer} free parameters for the blazar population than the more simplified model (three versus four free for the blazar model, plus those fixed in the nonblazar AGN model in this work). Most importantly, the Fermi-LAT analysis of FB10 fixes the spectral index of the blazar population, and, crucially, does not include the hardening of the spectra of the unresolved low-luminosity blazar population. The hardening of spectra with lower luminisity has been seen by both EGRET~\cite{Fossati97,*Fossati98,*Donato01} and Fermi-LAT (Fig.~\ref{point}). The fixed spectrum forces the FB10 conclusion that only $\sim$16\% of the GeV isotropic diffuse background could arise from blazars, and is also the case in other work using fixed blazar spectra \cite{Malyshev:2011zi}. Other recent work with different blazar population models, including spectral shape variation~\cite{Venters:2011gg}, possible point-source confusion~\cite{Stecker:2010di}, and BL Lac dominance of the unresolved portion~\cite{Neronov:2011kg} also find that a substantial portion of the DGRB could arise from the blazar population. Overall, the SED-sequence model of blazars and AGN as the source of the DGRB is remarkably consistent with the measured DGRB spectrum and blazar source-count distribution. The SED-sequence will continue to be improved with upcoming Fermi-LAT blazar data~\cite{Meyer:2011uk}. Further analyses of the type presented here, incorporating potential enhancements to the SED-sequence model, the XLF of AGN, and general studies of observed blazar spectral properties, will further enlighten the understanding of the extragalactic gamma-ray sky.

1012.1247

1012

1012.1071_arXiv.txt

We carried out extremely sensitive Submillimeter Array (SMA) 340 GHz continuum imaging on two submillimeter galaxies (SMGs): GOODS~850-11 and GOODS~850-13. The observations reach sub-mJy rms sensitivities and, interestingly, resolve both sources into multiple, physically unrelated SMGs. GOODS~850-11 is resolved into two sources at different redshifts. GOODS~850-13 is resolved into three sources, two with different spectroscopic redshifts and one only with a photometric redshift. All the SMA sources have fluxes in the 3--5~mJy range and all are detected at 1.4~GHz. Three of them are detected by \emph{Chandra}, and one is a previously unknown X-ray SMG. This is the first time that single-dish SMGs are resolved into multiple unrelated sources and also the first time that the SMA has discovered new SMGs. Our results show that identifications of SMGs at any wavelengths other than the submillimeter itself can be misleading, since such identifications usually only pick up one of the real counterparts. Using simulations that mimic our SCUBA and SMA observations, we find that the number of triple systems detected in our SMA survey is much higher than that expected from the current best-determined number counts. We tentatively attribute this to clustering. We also predict that ALMA will find $\sim1/3$ of $>5$~mJy 850 $\mu$m SCUBA sources to be multiple systems. Based on our SMA observations and simulations, we suggest that large samples of existing SMGs should be imaged with sensitive interferometric observations, even if the SMGs were previously thought to be securely identified.

Since the first discoveries of distant submillimeter galaxies (SMGs; \citealp{smail97,barger98,hughes98,eales99}), tremendous progress has been made in understanding their nature and their role in galaxy evolution. With single-dish submillimeter surveys, we are now able to resolve approximately 30\% of the 850~$\mu$m background into point sources with $S_{850~\mu \rm m}\gtrsim2$ mJy and constrain their number counts fairly well at this bright end \citep[e.g.,][]{coppin06}. However, the low resolution of single-dish telescopes and the associated large positional uncertainties make the identification and followup of these sources quite difficult. Roughly 60\%--70\% of bright SMGs have counterparts in deep radio interferometric images \citep{barger00,ivison02,chapman03b} and their positions are known with subarcsec accuracy. Spectroscopic followup of the radio identified SMGs shows a redshift distribution between $z\sim1.5$--3.5 and that the SMGs dominate the total star formation in this redshift range \citep{chapman03a,chapman05}. The recent advent of the Submillimeter Array \citep[SMA;][]{ho04} further helps to identify the radio-faint SMGs, and the redshift distribution has been extended to $z>4$ \citep{iono06,wang07,younger07,cowie09, younger09}. In addition to the redshift identifications, detailed followup observations have been made to study the properties of the SMGs in the X-ray, near-infrared, mid-infrared, and molecular line transitions (e.g., \citealp{alexander03a}, hereafter A03; \citealp{swinbank04,pope08,yun08,greve05,tacconi06}). The followup studies of SMGs have been overwhelmingly focused on the brighter sources ($S_{850~\mu \rm m}\gtrsim5$~mJy), which can be easily detected by single-dish telescopes. The nature of fainter SMGs that comprise the bulk of the submillimeter background has been much less explored. Surveys of lensing cluster fields yield small samples of sub-mJy sources \citep{blain99,cowie02,knudsen08}. The number counts indicate that the background is dominated by sources with $S_{850~\mu \rm m}\sim1$~mJy and that the full resolution of the background requires detections of $\sim0.1$~mJy sources. Stacking analyses statistically detect faint SMGs at 500 $\mu$m to 1.2~mm, and various attempts have been made to study the redshift distribution of these faint SMGs (e.g., Wang et al.\ 2006; \citealp{serjeant08,marsden09,penner10}). However, the results from these stacking analyses have not converged to a consistent picture. A full understanding of the submillimeter background population will likely require next-generation interferometers, such as the Atacama Large Millimeter/Submillimeter Array (ALMA). Fortunately, it is now possible to detect more typical submillimeter sources with the SMA. The recent upgrade of the SMA to a 4~GHz bandwidth not only greatly boosts its continuum sensitivity, but it also makes the calibrations with fainter quasars easier than before. This provides a first glance of what we may find in deep ALMA surveys and what issues may be present in current studies of SMGs. In this letter we report new SMA 340~GHz continuum observations of two SMGs, GOODS~850-11 and GOODS~850-13 (\citealp{wang04}, hereafter W04), aka.\ GN12 and GN21 \citep{pope05} in the Great Observatories Origins Deep Survey-North (GOODS-N; \citealp{giavalisco04}). Interestingly, both single-dish sources are resolved by the SMA into multiple physically unrelated galaxies. To our knowledge, these are the first examples of resolved, unrelated, multiple sources in the SMG population.

Since the commissioning of the SMA, it has been used for followup observations of SMGs, either for identification \citep[e.g.,][]{iono06,wang07,younger07,cowie09,younger09} or for morphology \citep[e.g.,][]{younger08}. This is the first time that the SMA has discovered new SMGs. The multiple SMA detections illustrate the limitations of identifying SMGs in any wavelength other than the submillimeter itself. Both sources had radio, 24 $\mu$m, and X-ray identifications in P06 and A03. All of the previously proposed identifications are only partially correct; i.e., they are all legitimate SMGs, but the submillimeter fluxes and source numbers will be misinterpreted by up to a factor of 3. If such cases are common, then our understanding of the SMG population is fundamentally flawed. The effects of multiple SMGs are commonly included in single-dish number counts simulations (e.g., \citealp{eales00,scott02,coppin06}; W04). However, its importance (relative to other effects such as Eddington bias) has not been directly demonstrated by observations. So far only sensitive millimeter/submillimeter interferometric observations can reveal the existence of multiple SMGs like GOODS~850-11 and 13, since existing single-dish telescopes are still severely confusion limited (at $>850~\mu$m) or noise limited (at 450~$\mu$m). The SMA surveys of \citet{younger07,younger09} imaged 15 SMGs selected at 1.1 mm. They did not find evidence of multiple sources, consistent with the argument made by \citet{ivison07} that multiple sources are rare. However, the SMA surveys of Younger et al.\ have 345 GHz rms sensitivities of 1--2 mJy. Even if there are secondary sources with $S_{\rm 345~GHz}\sim3$--5 mJy in their survey fields, such sources would not be easily detected. Our SMA survey is the first one that is deep enough to reveal such cases. After the SMA upgraded to the 4 GHz bandwidth, we observed five GOODS-N SMGs at similar depths (A.\ Barger et al., in preparation). Two of them are resolved into multiple sources and reported here. There is a third resolved source that will be reported by Barger et al. It is unclear whether this source is a merger or a physically unrelated pair. Even if we exclude the third source, the frequency of multiple sources in our SMA sample still seems unusually high. To better understand this, we performed Monte Carlo simulations that mimic our SCUBA and SMA surveys. We first created 1000 simulated SCUBA images using the differential counts in \citet{cowie02} and \citet{coppin06}, the ``true-noise'' map created in our GOODS-N SCUBA survey in W04, and the SCUBA beam in W04. We than extracted the simulated SCUBA sources. For each SCUBA source detected at $>5$ mJy and $>4~\sigma$ (the selection criteria for our SMA observations), we searched the input catalog for any $>3$ mJy (our SMA detection limit) sources within $17\arcsec$ (the SMA primary beam HWHM) of the SCUBA position. In the simulations the mean number of $>5$ mJy sources in a W04 SCUBA area is $11.4\pm3.9$, where the uncertainty is the dispersion in the 1000 realizations and is nearly Poissonian. This value is consistent with the W04 SCUBA observations (15 sources). On the other hand, it is significantly larger than the cumulative count, as it is affected by blending and flux boosting. The measured counts take these effects out. The mean numbers of double and triple systems are $1.29\pm1.33$ and $0.06\pm0.27$, respectively. We had observed 10 of the 15 $>5$ mJy SCUBA sources with the SMA. The most recent five of the 10 SMA observations were deep enough to detect $>3$ mJy sources. The earlier SMA observations were shallower, so there was a selection bias against SCUBA sources with multiple counterpart candidates. Given this bias and to be conservative, we only scale the above values of $1.29\pm1.33$ and $0.06\pm0.27$ by a factor of 10/15, rather than 5/15. We thus expect to find $0.86\pm0.89$ double and $0.04\pm0.18$ triple systems. The probability of finding one triple system like GOODS 580-13 is only 4\%, inconsistent with the actual observations. Among the possible explanations, the most likely one is clustering of SMGs, which is not included in the simulations. This is plausible because the photometric redshift of GOODS~850-13a has a confidence range (Table~\ref{tab1}) covering the redshift of 13b or 13c. This can be tested this with future spectroscopic observations in the near-infrared or in the millimeter. In the same simulations, we increased the SMA detection limit to 4~mJy, and we found that the probability of multiple systems dramatically decreases to $\sim6\%$ . This is consistent with the SMA survey results of \citet{younger07,younger09}. By altering the details of the simulations, we also found that the above results are fairly insensitive to the following SCUBA and SMA observing strategies: (1) the source extraction (various detection thresholds), (2) flux measurement in the SCUBA map (simple aperture flux vs.\ optimally filtered flux using the beam), (3) the shape of the SCUBA sidelobes (determined by the secondary chopping), and (4) the decision on where to point the SMA (as long as it is within the SCUBA positional uncertainty). We can use our simulations to predict the early results of ALMA identifications of SCUBA sources. We adopt the primary beam HWHM of $8\farcs5$ of ALMA at 340~GHz. In the ALMA early science phase we expect to have at least 10 antennae, and we can detect 0.5~mJy sources in roughly one hour. If we point ALMA at a $>5$~mJy SCUBA source, the probabilities of detecting double and triple SMGs within the primary beam will be 29\% and 6.5\%, respectively. The combined fraction of $\sim35\%$ is very high. We therefore predict that multiple detections in early ALMA observations will be quite common. At the beginning of this section, we raised the issue about incomplete identifications of SMGs in the X-ray, 24~$\mu$m, and radio. Based on our SMA observations and the above simulations, we believe that multiple systems and therefore incomplete identifications are common. Thus, we suggest that large numbers of single-dish sources should be re-identified with sensitive interferometric observations, even if the sources were previously thought to be securely identified.

1012.1071

1012

1012.3678_arXiv.txt

We present new theoretical estimates of the relative contributions of unresolved blazars and star-forming galaxies to the extragalactic \gray background (EGB) and discuss constraints on the contributions from alternative mechanisms such as dark matter annihilation and truly diffuse \gray production. We find that the {\it Fermi} source count data do not rule out a scenario in which the EGB is dominated by emission from unresolved blazars, though unresolved star-forming galaxies may also contribute significantly to the background, within order-of-magnitude uncertainties. In addition, we find that the spectrum of the unresolved star-forming galaxy contribution cannot explain the EGB spectrum found by EGRET at energies between $50$ and $200$ MeV, whereas the spectrum of unresolved FSRQs, when accounting for the energy-dependent effects of source confusion, could be consistent with the combined spectrum of the low-energy EGRET EGB measurements and the {\it Fermi}-LAT EGB measurements.

Studies of the extragalactic \gray background (EGB) can provide insight into high energy processes in the universe and, as such, has been the subject of much debate, particularly concerning the roles of extragalactic astrophysical sources and new physics. Recent data from the Large Area Telescope (LAT)\footnote{Hereafter we shall refer to the {\it Fermi}-LAT instrument simply as {\it Fermi}.} on board the {\it Fermi Gamma Ray Space Telescope} allow for a reassessment of the possible astrophysical origins of the EGB, which could improve our understanding of \gray production in these objects and provide more robust constraints on the more exotic scenarios. However, in order to determine the strength and spectrum of this isotropic background one needs to proceed from the raw photon count data by determining to the best extent possible the detector sensitivity, the intrinsic events produced by the larger charged particle flux impinging on the detector, and the much larger \gray foreground within our Galaxy resulting from cosmic-ray interactions with photons and gas nuclei. Such analyses have been made for both the Energetic \gray Experiment Telescope (EGRET) aboard the {\it Compton Gamma Ray Observatory} \citep{sre98,str04} and {\it Fermi} \citep{lat10}. Various extragalactic \gray production scenarios have been explored theoretically as candidate components that could contribute significantly to the observed background. Among those considered are the unresolved astronomical sources, such as active galactic nuclei (AGN) \citep{pad93,ste93,sal94,chi95,ste96,kaz97,chi98,muk99,muc00,gio06,nt06,der07,pv08,ino09,ven09, ven10,aba10}, star-forming galaxies \citep{pav02, fie10, mak10}, and starburst galaxies \citep{tho07, ste06, mak10}. The large majority of associated extragalactic sources thus far detected by both EGRET and {\it Fermi} are blazars \citep{har99,1cat}, {\it i.e.}, those AGN for which the jet is closely aligned with the observer's line-of-sight \citep{bla79}, including \gray loud flat spectrum radio quasars (FSRQs) and BL Lacertae-type objects. It is expected that since blazars comprise the largest class of identified extragalactic \gray sources, unresolved blazars should contribute significantly to the EGB. Additionally, just as our Galaxy produces $\gamma$-rays, it is expected that \grays are produced in other galaxies, and as such, unresolved galaxies might also contribute to the EGB with the most significant contribution originating from the population of actively star-forming galaxies \citep{ste75, pav01,pav02,fie10,mak10}. Interesting truly diffuse mechanisms that could contribute to the EGB involve cosmic ray interactions with intergalactic gas and the cosmic background radiation \citep{faz66,ste73,dar07,kes03} and electromagnetic cascades produced by interactions of very high and ultrahigh energy particles with the extragalactic background light \citep{kal09, ber10, ahl10,ven10}, as well as more exotic scenarios such as dark matter annihilation~\citep{sil84, ste85, rud88, ste89, ste89a, rud91,ull02} and decay \citep{oli85,ste86, iba08}\footnote{For reviews on dark matter annihilation, see \citet{jun96} and \citet{ber05}.}. In this paper, we estimate the contributions to the EGB from unresolved extragalactic \gray sources of various types and compare them with the EGB obtained from analysis of {\it Fermi} data. In doing so, we also take into consideration the effects of both the completeness of the {\it Fermi} flux limited blazar survey and the important effect of source confusion owing to the energy dependent angular resolution of the {\it Fermi}-LAT detector. We will then briefly discuss the implications of possible truly diffuse emission mechanisms to the EGB.

We have calculated the spectral shape of the contribution of unresolved FSRQs to the EGB assuming that the \gray luminosity of an FSRQ is, on average, proportional to its radio luminosity \citep{gir10, lat10,ghi10a,mah10}, and also accounting for the effects of source confusion. We have demonstrated that the combination of the source density predicted by the \citet{dun90} FSRQ radio luminosity function and the strong energy dependence of the {\it Fermi}-LAT angular resolution \emph{increases} the contribution of unresolved FSRQs to the EGB at energies below $1$ GeV. The resulting overall spectrum predicted by the fit to the {\it Fermi} source count distribution reproduces well the spectrum of the EGRET and {\it Fermi} EGB measurements below $1$ GeV, but falls below the data points above $1$ GeV. We have also calculated the spectral shape of the contribution of unresolved star-forming galaxies to the EGB for several relations for the \gray luminosity of a star-forming galaxy. We find that, depending on the model, the overall amount of the contribution of star-forming galaxies to the EGB may be more or less significant, though regardless of the model considered, the spectrum of unresolved star-forming galaxies is unable to explain the combined spectrum of the low-energy EGRET EGB measurements and the {\it Fermi} EGB measurements. Similar calculations for starburst galaxies alone indicate that they account for at most about $1$\% of the EGB, in agreement with the conclusion reached by \citet{ste06}. The similarity of the collective spectrum of unresolved FSRQs to the combined spectrum of the EGRET and {\it Fermi} EGB measurements, as demonstrated by our results, is striking. In fact, we note that as predicted in \citet{ste99}, the inclusion of the effect of source confusion in the calculation could provide an explanation for the similarity between the EGRET and {\it Fermi} EGB measurements at energies of hundreds of MeV. The density of FSRQs predicted by the model is sufficiently large such that at these energies, {\it Fermi} would not be able to resolve many more FSRQs than EGRET did, and the FSRQ contribution to the EGB would remain the same for {\it Fermi} as for EGRET. Thus, if unresolved FSRQs \emph{do} comprise the bulk of the EGB emission, then one would expect such similarity between the EGRET and {\it Fermi} measurements at these energies. At energies above $1$ GeV, the {\it Fermi}-LAT angular resolution improves substantially with respect to that of EGRET, and as such, {\it Fermi} would be able to resolve more blazars at higher energies than EGRET could, resulting in a decrease in the {\it Fermi} EGB with respect to the EGRET EGB, an effect which is possibly indicated by comparing the EGRET and {\it Fermi} results\footnote{A caveat is that the uncertainties in the subtraction of the galactic foreground emission at the higher energies are considerable owing to the uncertainty in the distributions of both gas and cosmic rays in the Galaxy. Furthermore, the instrumental backgrounds of EGRET and the {\it Fermi}-LAT are different, so it is difficult to make a direct comparison between the two. We should also note that in our calculations, we have neglected the population of BL Lacs, which, due to their hard spectra, are likely to have more of a contribution at energies above $\sim 10$ GeV. Notably, {\it Fermi} has resolved as many BL Lacs as FSRQs.}. In contrast, no such high-energy separation between the EGRET EGB and the {\it Fermi} EGB is predicted for star-forming galaxies as all but the closest are too faint to be resolvable by {\it Fermi}. Furthermore, we note that the EGRET EGB measurements provide no indication of a turnover in the spectrum as would be expected if unresolved star-forming galaxies comprise the bulk of the EGB emission. Rather, the spectrum of unresolved star-forming galaxies is \emph{inconsistent} with the combined spectrum of the EGRET and {\it Fermi} EGB measurements\footnote{As previously noted, the {\it Fermi}-LAT was designed to reach its optimal effective area for \grays with energies near and above $\sim 1$ GeV, whereas EGRET was designed to reach its optimal effective area for \grays with energies near and above $\sim 100$ MeV. Also, the {\it Fermi}-LAT detector has a significantly higher instrumental background at $100$ MeV than EGRET did (S. D. Hunter, private communication). Thus, the EGB was not reported by {\it Fermi} for energies below $200$ MeV \citep{lat10}.}. We also note that the lack of a turnover in the EGRET data is not simply the result of systematics (S. D. Hunter, private communication), since the uncertainties in all of the galactic foreground models used to determine the EGB from the EGRET and {\it Fermi} data are quite small at these energies. Finally, we note that at energies above $\sim 1$ GeV, the spectrum of unresolved star-forming galaxies is steeper than the spectra of the EGB data. As such, we conclude that however significant the contribution of star-forming galaxies to the EGB may be, it is not sufficient to explain the EGB\footnote{The effect of Compton interactions mentioned in Section \ref{subsec:galaxyform} does not alter this conclusion as it only modifies the spectrum above 10 GeV for normal galaxies \citep{str10} and the starburst galaxy contribution to the EGB is negligible (See Figure \ref{fig:egrbsfgal}.)}. Within the range of our various predictions of the EGB from star forming galaxies, we agree with the results of the models of both~\cite{fie10} (which suggests that star-forming galaxies may comprise the bulk of the EGB) and \citet{mak10} (which suggests that star-forming galaxies can account for less than $10\%$ of the EGB). This underscores the range of uncertainty in the calculation for star-forming galaxies\footnote{One noteworthy difference between our model and that of \citet{fie10} is that in order to relate the gas mass of a galaxy to its star formation rate, \citet{fie10} makes use of the Schmidt-Kennicutt relation. However, in doing so, they estimate the disk sizes of galaxies to high redshifts. Given that the uncertainties in these quantities are likely to be considerable (and given the uncertainty already present in the calculation), we considered alternative approaches. Nevertheless, we note that the Schmidt-Kennicutt relation was included in the inputs to both the IR luminosity models and the Schechter function model. In the IR luminosity model, we tested the impact of changing the Schmidt-Kennicutt law by performing the calculation for the \citet{hop10} IR luminosity function calculated using a steeper Schmidt-Kennicutt relation and found that it had very little impact on our results.}. The featureless spectrum of the EGB deduced by {\it Fermi} is intriguing when one considers the possibility of features that could arise from phenomena such as breaks in blazar spectra, absorption of high-energy \grays from unresolved blazars, \gray emission from unresolved star-forming galaxies, \gray emission from dark matter annihilation, and \grays from electromagnetic cascades initiated by very high and ultrahigh energy particle interactions with the extragalactic background light. The spectra of these potential contributions to the EGB differ considerably from that of the FSRQs \citep{sil84,ste85,rud88,ste89, ste89a,rud91,ull02,ando07,kal09,sie09,ber10,ahl10,ven10}. However, recent {\it Fermi} observations have placed significant constraints on dark matter annihilation \citep{cir09,dm10, ack10}, and presently there is no clear evidence of annihilation features above the background continuum. As such, it appears that any putative contribution to the EGB from dark matter annihilation is relatively minor. The possible contribution to the EGB from electromagnetic cascades is constrained by the relative steepness of the EGB spectrum, though cascades could play a role at higher energies \citep{kal09,ber10,ahl10,ven10}. An apparent explanation for the featureless power-law spectrum of the EGB as presently deduced could be that unresolved blazars provide the dominant contribution to the EGB, given that their collective spectrum is roughly consistent with that of the EGB. Therefore, we conclude that, contrary to the result given by~\citet{pop10}, the {\it Fermi} observations do not rule out the possibility that the EGB is dominated by emission from unresolved blazars.

1012.3678

1012

1012.4869_arXiv.txt

When one splits spacetime into space plus time, the spacetime curvature (Weyl tensor) gets split into an ``electric'' part $\mathcal E_{jk}$ that describes tidal gravity and a ``magnetic'' part $\mathcal B_{jk}$ that describes differential dragging of inertial frames. We introduce tools for visualizing $\mathcal B_{jk}$ (frame-drag vortex lines, their vorticity, and vortexes) and $\mathcal E_{jk}$ (tidal tendex lines, their tendicity, and tendexes), and also visualizations of a black-hole horizon's (scalar) vorticity and tendicity. We use these tools to elucidate the nonlinear dynamics of curved spacetime in merging black-hole binaries.

1012.4869

1012

1012.3164_arXiv.txt

{ A milestone of modern cosmology was the prediction and serendipitous discovery of the Cosmic Microwave Background (CMB), the radiation left over after decoupling from matter in the early evolutionary stages of the Universe. A prediction of the standard hot Big-Bang model is the linear increase with redshift of the black-body temperature of the CMB ($T_{\rm CMB}$). This radiation excites the rotational levels of some interstellar molecules, including carbon monoxide (CO), which can serve as cosmic thermometers. Using three new and two previously reported CO absorption-line systems detected in quasar spectra during a systematic survey carried out using VLT/UVES, we constrain the evolution of $T_{\rm CMB}$ to $z\sim3$. Combining our precise measurements with previous constraints, we obtain $T_{\rm CMB}(z)=(2.725\pm0.002)\times (1+z)^{1-\beta}$~K with $\beta=-0.007\pm0.027$, a more than two-fold improvement in precision. The measurements are consistent with the standard (i.e. adiabatic, $\beta=0$) Big-Bang model and provide a strong constraint on the effective equation of state of decaying dark energy (i.e. $w_{eff}=-0.996\pm0.025$).}

The existence of the Cosmic Microwave Background (CMB) radiation is a fundamental prediction of the hot Big-Bang theory. If gravitation is described by general relativity and electromagnetism by Maxwell theory then photons propagate along null geodesics and the CMB black-body temperature must follow the relation $T_{\rm CMB}(z)=T_{\rm CMB}^0\times(1+z)^{1-\beta}$, with $\beta=0$ and where $T_{\rm CMB}^0$~=~2.725$\pm$0.002~K \citep{Mather99} is the temperature measured locally (at redshift $z=0$). This relation, which is a theoretical consequence of the adiabatic expansion of the Universe, needs to be verified by direct measurements. This has also deeper theoretical implications \citep[][]{Uzan04}. A non-zero $\beta$ would indicate either a violation of the hypothesis of local position invariance (and thus of the equivalence principle) or that the number of photons is not conserved -- with the constraint that the energy injection does not induce spectral distortion of the CMB. In the first case, this should be associated with a variation of the fundamental constants \citep[see e.g.][]{Murphy03,Srianand04}. There are currently two methods to measure $T_{\rm CMB}$ at redshifts $z>0$. The first one relies on the measurement of a small change in the spectral intensity of the CMB toward clusters of galaxies due to inverse Compton scattering of photons by the hot intra-cluster gas: the so-called Sunyaev-Zel'dovich (S-Z) effect \citep{Fabbri78,Rephaeli80}. Although this technique permits precise measurements \citep[$\Delta T \sim 0.3$\,K;][]{Battistelli02,Luzzi09}, the method is essentially limited to $z<0.6$ because of the scarcity of known clusters at higher redshifts. The other technique uses the excitation of interstellar atomic or molecular species that have transition energies in the sub-millimetre range and can be excited by CMB photons. When the relative population of the different energy levels are in radiative equilibrium with the CMB radiation, the excitation temperature of the species equals that of the black-body radiation at that redshift. Therefore, the detection of these species in diffuse gas, where collisional excitation is negligible, provides one of the best thermometers for determining the black-body temperature of the CMB in the distant Universe \citep[][]{Bahcall68}.

\begin{figure*} \centering \includegraphics[bb=24 124 566 396,clip=, width=0.9\hsize]{figtcmb_coSZ.ps} \caption{The black-body temperature of the Cosmic Microwave Background radiation as a function of redshift. The star represents the measurement at $z=0$ \citep{Mather99}. Our measurements based on the rotational excitation of CO molecules are represented by red filled circles at $1.7<z<2.7$. Other measurements at $z>0$ are based (i) on the S-Z effect (blue triangles at $z<0.6$, \citealt{Luzzi09}) and (ii) on the analysis of the fine structure of atomic carbon (green open squares: $z=1.8$, \citealt{Cui05}; $z=2.0$, \citealt{Ge97}; $z=2.3$, \citealt{Srianand00}; $z=3.0$, \citealt{Molaro02}). Upper-limits come from the analysis of atomic carbon (from the literature and our UVES sample, see \citealt{Srianand08}) and from the analysis of molecular absorption lines in the lensing galaxy of PKS\,1830-211 \citep[open circle at $z=0.9$,][]{Wiklind96}. The dotted line represents the adiabatic evolution of $T_{\rm CMB}$ as expected in standard hot Big-Bang models. The solid line with shadowed errors is the fit using all the data and the alternative scaling of $T_{\rm CMB}(z)$ \citep{Lima00} yielding $\beta=-0.007\pm0.027$. The red dashed curve (resp. green dashed-dotted) represents the fit and errors using S-Z + CO measurements (resp. S-Z + atomic carbon). \label{TCMBz}} \end{figure*} The CMB temperature derived from the rotational excitation of CO in five absorption systems (three studied here plus two previously published, see Table~\ref{tco}) are presented in Fig.~\ref{TCMBz} together with measurements and upper-limits obtained from the analysis of the populations of the fine-structure energy levels of atomic carbon or from the S-Z effect in galaxy clusters. An upper-limit has also been obtained from the analysis of millimetre absorption lines from different molecules in the gravitational lens of PKS\,1830-211 \citep{Wiklind96}. The technique presented here allows us to probe the temperature of the CMB at high redshift, providing constraints that are independent from and stronger than those arising from the analysis of C$^0$ and C$^+$. This demonstrates that the rotational excitation of interstellar CO molecules can provide a direct and precise measurement of $T_{\rm CMB}$ in the early Universe. Fitting the measurements from different techniques with the expression $T_{\rm CMB}(z)=T_{\rm CMB}^0\times(1+z)^{1-\beta}$ \citep{Lima00}, we get the constraints on $\beta$ summarised in Table~\ref{beta}. We note that combining CO measurements of $T_{\rm CMB}(z)$ with those obtained from other techniques, improves the precision of the $\beta$-measurement by more than a factor of two. The measurement presented here, $\beta=-0.007\pm0.027$, directly supports the adiabatic evolution of the CMB radiation temperature ($\beta=0$), expected from the standard hot Big-Bang model. Considering alternative $\Lambda$ cosmological models, \citet{Jetzer10} demonstrated that measuring $T_{\rm CMB}$ at different redshifts allows one to constrain the effective equation of state of decaying dark energy ($p=w_{eff} \rho$). Fitting the measurements of $T_{\rm CMB}$ with their temperature-redshift relation (Eq. 22 in \citealt{Jetzer10}), taking $\Omega_m=0.275\pm0.015$ \citep{Komatsu10} and fixing $\gamma$ to the canonical value (4/3), we get $w_{eff}=-0.996\pm0.025$ which is a tighter constraint compared to those previously derived from other methods \citep[e.g.][]{Kowalski08,Riess09,Kessler09,Jullo10}. Finally, large and deep QSO surveys such as SDSS~III should provide more lines of sight along which CO can be detected while high-resolution spectrographs on future extremely large telescopes will allow for full de-blending of the CO lines in different rotational levels, yielding more accurate measurements. \begin{table} \caption{Constraints on the evolution of $T_{\rm CMB}$ with redshift \label{beta}} \centering \begin{tabular}{cc} \hline \hline Data set & $\beta$ \\ \hline S-Z & $+0.040\pm0.079$ \\ S-Z + atom. carbon & $+0.029\pm0.053$ \\ S-Z + CO & $-0.012\pm0.029$ \\ S-Z + atom. carbon + CO & $-0.007\pm0.027$ \\ \hline \end{tabular} \end{table}

1012.3164

1012

1012.3214_arXiv.txt

We present new results on the "dark flow" from a measurement of the dipole in the distribution of peculiar velocities of galaxy clusters, applying the methodology proposed and developed by us earlier. Our latest measurement is conducted using new, low-noise 7-yr WMAP data as well as an all-sky sample of X-ray selected galaxy clusters compiled exclusively from published catalogs. Our analysis of the CMB signature of the kinematic Sunyaev-Zeldovich (SZ) effect finds a statistically significant dipole at the location of galaxy clusters. The residual dipole outside the cluster regions is small, rendering our overall measurement 3-4 sigma significant. The amplitude of the dipole correlates with cluster properties, being larger for the most X-ray luminous clusters, as required if the signal is produced by the SZ effect. Since it is measured at zero monopole, the dipole can not be due to the thermal SZ effect. Our results are consistent with those obtained earlier by us from 5-yr WMAP data and using a proprietary cluster catalog. In addition, they are robust to quadrupole removal, demonstrating that quadrupole leakage contributes negligibly to the signal. The lower noise of the 7-yr WMAP also allows us, for the first time, to obtain tentative empirical confirmation of our earlier conjecture that the adopted filtering alters the sign of the KSZ effect for realistic clusters and thus of the deduced direction of the flow. The latter is consistent with our earlier measurement in both the amplitude and direction. Assuming the filtering indeed alters the sign of the KSZ effect from the clusters, the direction agrees well also with the results of independent work using galaxies as tracers at lower distances. We make all maps and cluster templates derived by us from public data available to the scientific community to allow independent tests of our method and findings.

Peculiar velocities play an important role in understanding the large-scale gravitational field in the Universe and have been the subject of intense investigations over the past decades. Early determinations of peculiar velocities were based on surveys of individual galaxies (see review by Strauss \& Willick 1995). First measurements by Rubin and co-workers found peculiar flows of ${\sim}700$ km s$^{-1}$ (Rubin et al 1976), but were largely dismissed at the time. A group collectively known as the ``Seven Samurai" found that elliptical galaxies within ${\sim}60h^{-1}$Mpc were streaming at ${\sim}600$ km s$^{-1}$ with respect to the CMB (Dressler et al 1987, Lynden-Bell et al 1988). Using mainly spiral galaxies, Mathewson et al (1992) found that this flow does not converge until scales much larger than ${\sim}60 h^{-1}$ Mpc, in agreement with the results of a later analysis by Willick (1999). With brightest cluster galaxies as distance indicators for a sample of 119 rich clusters, Lauer \& Postman (1994 - LP) measured a bulk flow of ${\sim}$700 km s$^{-1}$ on a scale of ${\sim}150h^{-1}$Mpc. An improved re-analysis of these data by Hudson \& Ebeling (1997), however, found a reduced bulk flow pointing in a different direction. Using early-type galaxies in 56 clusters, Hudson et al (1999) found a similar bulk flow as LP and on a comparable scale, but again in a different direction. By contrast, a sample of 24 SNIa by Riess et al (1997) showed no evidence of significant bulk flows out to ${\sim}100 h^{-1}$ Mpc, and a similar conclusion was reached in a study of spiral galaxies by Courteau et al (2004). A complementary technique aimed at constraining bulk motions reconstructs directly the peculiar gravity of the observed galaxy distribution and uses measurements of the dipole in the distributions of light and matter. The dipole derived from the distribution of galaxies mapped in optical surveys is nearly aligned with the one obtained for infra-red-selected galaxies, but both are misaligned with respect to the CMB dipole generated by the motion of our Local Group relative to the CMB rest frame (Rowan-Robinson et al 2000), although this misalignment becomes less troublesome if one relaxes the light-tracing-mass assumptions (see discussion by Gunn 1988). Kocevski et al (2004) and Kocevski \& Ebeling (2006) measured the dipole anisotropy of an all-sky sample of X-ray-selected galaxy clusters to probe mass concentrations beyond the Great Attractor and found that most of the peculiar velocity of the Local Group is due to overdensities at $\ga 150h^{-1}$Mpc. All galaxy techniques lose sensitivity at distances approaching and greater than $\sim$50-100 Mpc. The Sunyaev-Zel'dovich (SZ) effect, produced by hot gas in galaxy clusters, is uniquely suited to probe flows to larger distances; moreover, it is independent of redshift and not subject to the systematics plaguing studies using empirical distance indicators. The kinematic part of the SZ effect (KSZ) is directly proportional to the cluster velocity with respect to the cosmic micorwave background (CMB). Because of its smallness, the KSZ effect has, however, not yet been measured for individual clusters; observations of six clusters at a wide range of redshifts out to $z\simeq 0.82$ yielded an upper limit of $V\la 1,500$ km s$^{-1}$ on a poorly defined scale (Benson et al 2003). Kashlinsky \& Atrio-Barandela (2000, hereafter KA-B) have proposed a method to measure large-scale flows using all-sky cluster catalog and CMB all-sky data, such as obtained with WMAP. KA-B identified a statistic (the dipole of the CMB temperature field evaluated at cluster positions) which preserves the KSZ component while integrating down other (noise) terms. However, the method requires a CMB filter that removes the primary CMB (which is strongly spatially correlated) without significantly attenuating the KSZ bulk flow contribution; clearly not every filter will achieve this. Kashlinsky et al (2008, 2009 - KABKE1,2) have applied the KA-B method to a large cluster catalog finding a surprising flow (dubbed the "dark flow") extending to at least 300$h^{-1}$Mpc. Following this, an independent study of Watkins et al (2010) combined the available galaxy data suppressing the sampling noise in the various surveys and showed that all data (with the exception of the LP sample) agreed with a substantial motion on a scale of ${\simeq}50-100 h^{-1}$Mpc with amplitude and direction in good agreement with the KABKE measurements. In a follow-up study, Kashlinsky et al (2010, KAEEK) revise the statistical analysis of their original study\footnote{Keisler (2009) together with KAEEK and AKEKE pointed out that KABKE did not account for correlations of the residual primary CMB fluctuations in the 8 DA channels of the WMAP data used in their analysis.} and use a much expanded cluster catalog, binned by cluster X-ray luminosity ($L_X$) to demonstrate that the CMB dipole increases with the $L_X$-threshold as required by the KSZ origin of the signal; such an $L_X$ dependence of the dipole is inconsistent with it originating from some putative systematic effect from primary CMB fluctuations. KAEEK find that the "dark flow" flow extends to at least $\ga 800$ Mpc, twice the distance reported by KABKE. Atrio-Barandela et al (2010, AKEKE), developed a formalism to understand -- both analytically and numerically -- the uncertainties in measurements using the KABKE filter; the same formalism is applicable to any filtering scheme. In addition, AKEKE demonstrate that the KABKE filter removes primary CMB fluctuations down to the fundamental limit of cosmic variance, rendering it optimal for such studies. Very recently, the dark flow results of KABKE/KAEEK have found support by a study (Ma et al 2010) using a compilation of galaxy distance indicators which reports the same "tilt" velocity as the dark flow and pointing in the same direction, within the calibration uncertainties discussed in KABKE2/KAEEK/AKEKE. On the other hand, the KABKE results have been challenged by Keisler (2009). Replicating the analysis of KABKE1,2 using a cluster catalog compiled from publicly available data, Keisler confirmed the central dipole values measured by KABKE2, but claimed that it has only marginal statistical significance. AKEKE (Sec.~4 and Fig.~5) have since shown that Keisler's error estimates are erroneous\footnote{Clearly correlations between 8 DA channels used in the studies can at most increase the KABKE1,2 errors by $\sqrt{8}$, whereas Keisler claimed a $\ga\sqrt{20}$ increase.} and largely due to him not having removed the monopole and dipole from the CMB maps {\it outside} the mask. In a more recent challenge Osborne et al.\ (2010) have likewise used publicly available X-ray cluster data, applied alternative filtering schemes and claimed not to be able to replicate the "dark flow" results.

This paper demonstrates - using public X-ray data - the existence of a statistically significant dipole associated exclusively with clusters. The dipole signal is highly statistically significant and remains at apertures containing zero monopole. Its amplitude further increases with the X-ray luminosity threshold of the cluster subsamples as it should if produced by the SZ terms. However, the fact that it arises at zero monopole precludes any significant TSZ contributions to the signal as discussed in KABKE2 and AKEKE. We believe that the only explanation of this measurement is a large-scale bulk flow. Any alternative explanation of the signal has so far not been suggested in the literature. Adopting the large-scale-flow interpretation of the measurement, the properties of the flow (amplitude, direction and variation with depth) are fully consistent with Fig. 2 and Table 1 of AKEKE. Adopting the bulk-flow interpretation of the measured dipole, with the calibration coefficients for this configuration from Table 1 of KAEEK, the flow amplitude would be $\sim$1,000 km s$^{-1}$ in the direction given by Eq. \ref{eq:dipole}. The amplitude and the direction of the flow are consistent with being constant at depths $zcH_0 \sim 300-550\,h^ {-1}$ Mpc. Note however the caveat that systematic calibration uncertainties likely cause us to overestimate the amplitude by up to 30\% (KABKE2), and that interpreting the direction of the flow from the KSZ effect in the filtered maps remains subject to a sign change for which we present first tentative empirical evidence in Fig.\ref{fig:f5}. In order to better probe and expand on our earlier ``dark flow" results we have designed an experiment named SCOUT (SZ Cluster Observations as probes of the Universe's Tilt). The SCOUT goals are to compile a sample of $\sim 1,500$ X-ray selected galaxy clusters with spectroscopically measured redshifts out to significantly greater distances than the current $z{=}0.25$ limit, and to apply the KA-B method to the 9-year WMAP and Planck maps. The latter mission, with its low noise, higher angular resolution and wider frequency coverage, will be particularly useful in calibrating the measurements. First SCOUT results from a preliminary sample of $\sim 1,000$ clusters have been reported in KAEEK. While the SCOUT catalog is being assembled, we have shown in this paper that the basic dark flow results can be readily verified using publicly available cluster data. We make the sample generated from this database available upon publication at \url{http://www.kashlinsky.info/bulkflows/data\_public} and encourage the community to test our findings using the tools provided there. In addition, this paper further addresses an important calibration issue resulting from our filtering of the CMB maps. Using 7-year WMAP W-channel data, we show empirically that filtering may lead to a sign change in the KSZ term (see KAEEK). To resolve this issue and improve the calibration we need to decrease the noise in the measurement and properly recalibrate the catalog of cluster properties; both of these goals are achievable with a larger SCOUT sample. Application of our method to Planck data in the 217 GHz channel, proposed by us earlier, will then allow accurate measurements of the velocity and direction of the flow. We acknowledge NASA NNG04G089G/09-ADP09-0050 and FIS2009-07238/GR-234/SyEC CSD 2007-00050 grants from Spanish Ministerio de Educaci\'on y Ciencia/Junta de Castilla y Le\'on. We thank our collaborators on the SCOUT/"dark flow" project, Dale Kocevski and Alastair Edge, for their numerous valuable contributions to the project. \clearpage

1012.3214

1012

1012.0782.txt

We present the goals and preliminary results of an unbiased, near-infrared, narrow-band imaging survey of the First Galactic Quadrant (10\deg\,$< l <$\,65\deg\ ; $-1.3$\deg\,$< b <$\,+1.3\deg). This area includes most of the Giant Molecular Clouds and massive star forming regions in the northern hemisphere. The survey is centred on the 1-0\,S(1) ro-vibrational line of \htwo, a proven tracer of hot, dense molecular gas in star-forming regions, around evolved stars, and in supernova remnants. The observations complement existing and upcoming photometric surveys (Spitzer-GLIMPSE, UKIDSS-GPS, JCMT-JPS, AKARI, Herschel Hi-GAL, etc.), though we probe a dynamically active component of star formation not covered by these broad-band surveys. Our narrow-band survey is currently more than 60\,\% complete. The median seeing in our images is 0.73\arcsec. The images have a 5\,$\sigma$ detection limit of point sources of K$\sim$18\,mag and the surface brightness limit is 10$^{-19}$\,W\,m$^{-2}$\,arcsec$^{-2}$ when averaged over our typical seeing. Jets and outflows from both low and high mass Young Stellar Objects are revealed, as are new Planetary Nebulae and - via a comparison with earlier K-band observations acquired as part of the UKIDSS GPS - numerous variable stars. With their superior spatial resolution, the UWISH2 data also have the potential to reveal the true nature of many of the Extended Green Objects found in the GLIMPSE survey.

Feedback from star formation, particularly massive star formation, has a radical impact on the nature of the interstellar medium (ISM) in galaxies. Outflows from protostars and radiated energy from high-mass young stars heat, excite, modify the chemistry of and may provide the turbulent motions in Giant Molecular Clouds (GMCs). Ultimately, massive stars also enhance metal abundances in the ISM. Understanding the formation of stars, particularly massive stars, is thus of crucial importance. To help us better understand the dynamical processes associated with massive star formation, but also to search for other line-emitters (Supernova remnants, Planetary nebulae, etc.) along the Milky Way, we are conducting an unbiased survey of the Spitzer Space Telescope GLIMPSE-North portion of the Galactic Plane in \htwo\ 1-0\,S(1) emission at 2.122\,$\mu$m. \htwo\ observations highlight regions of shocked or fluorescently excited molecular gas (T $\approx$ 2000\,K, n$_{\rm H_2} >$\,10$^3$\,cm$^{-3}$) and thus trace outflows and jets from embedded young stars, but also the radiatively excited boundary regions between massive stars and the ISM. Our unique survey -- dubbed the UKIRT Wide Field Infrared Survey for \htwo\, or UWISH2 -- complements the existing UKIRT Infrared Deep Sky Survey (UKIDSS) of the Galactic Plane (the UKIDSS GPS; Lawrence et al. \cite{2007MNRAS.379.1599L}, Lucas et al. \cite{2008MNRAS.391..136L}), as well as the Isaac Newton Telescope Photometric H$\alpha$ Survey (IPHAS: Drew et al. \cite{2005MNRAS.362..753D}), the Herschel Infrared Galactic Plane Survey (Hi-GAL: Molinari et al. \cite{2010PASP..122..314M}), the planned James Clerk Maxwell Telescope Submillimeter Common User Bolometer Array-2 (SCUBA-2) Galactic Plane Survey (the JPS), the Very Large Array 5\,GHz "CORNISH" survey (Purcell et al. \cite{2008ASPC..387..389P}), the AKARI mid-infrared all sky survey (Ishihara et al. \cite{2010A&A...514A...1I}), and of course the mid-infrared GLIMPSE survey with Spitzer (Benjamin et al. \cite{2003PASP..115..953B}, Churchwell et al. \cite{2006ApJ...649..759C}) as well as the MIPSGAL survey (Carey et al. \cite{2009PASP..121...76C}). Other surveys which are covering a large fraction of our field are the Bolocam Northern Galactic Plane Survey at 1.1\,mm (BGPS, Rosolowsky et al. \cite{2010ApJS..188..123R}, Aguirre et al. \cite{2010arXiv1011.0691A}), the APEX Telescope Large Area Survey of the Galaxy at 870\,$\mu$m (ATLASGAL, Schuller et al. \cite{2009A&A...504..415S}), the Millimetre Astronomy Legacy Team 90 GHz Survey (MALT\,90), the Galactic Ring Survey (GRS, Jackson et al. \cite{2006ApJS..163..145J}) and the Methanol Multibeam Survey (MMB, Green et al. \cite{2007IAUS..242..218G}). Together the GLIMPSE, UKIDSS-GPS, JPS, Hi-GAL, CORNISH and other surveys provide a near-complete census of star formation, detecting cool, pre-stellar cores, hot cores, low and high-mass protostars (Class\,0/I), pre-main sequence objects (Class\,II/III -- T\,Tauri and Herbig Ae/Be stars) and H{\small II} regions associated with intermediate and high-mass young stars. The longer wavelength surveys will also map the associated dust distribution. Outflows from low and high-mass protostars, as well as H{\small II} regions, Photo-Dissociation Regions (PDRs), post-Asymptotic Giant Branch (post-AGB) stars and planetary nebulae (PN) are detectable by the UWISH2 survey and, in many cases, are resolved out to distances of at least 5\,kpc. \begin{figure*} \includegraphics[width=6cm]{UWISH2-Fig1.pdf} \\ \includegraphics[width=6cm]{UWISH2-Fig2.pdf} \caption[]{\label{area} {\bf Left:} An overview of the UWISH2 target region in the Galactic Plane. The image shows the IRAS dust emission at 100\mic\ with the Equatorial coordinate system overlayed in green and the target region in white. {\bf Right:} A map showing the completeness of the UWISH2 survey observations as of 21 September 2010 (i.e. after our observing run in Semester 2010B). Blue symbols mark observed tiles; green symbols mark remaining tiles. As of September 2010 the survey is 62.9\,\% completed. The symbols mark the four positions observed within each tile. Note that there are no gaps between adjacent tiles.} \end{figure*}

The UWISH2 narrow-band imaging survey is being used to trace dynamic processes associated with star formation and late stellar evolution. In particular, it picks out active regions of star formation, leading to estimates of star formation efficiency along the Galactic Plane. At the same time it also yields a more complete, unbiased census of post-AGB stars and PN in the Milky Way, leading to a comprehensive catalogue of PN morphologies as well as a map of their distribution along the Inner Galaxy. A major strength of the UWISH2 survey lies in its complementarity with other existing or upcoming surveys; the benefits of combining high-spatial resolution narrow-band WFCAM images of massive star forming regions with mid- and far-IR observations has recently been demonstrated by Davis et al. \cite{2007MNRAS.374...29D} and Kumar et al. \cite{2007MNRAS.374...54K}. In these complex environments the narrow-band data yield flow statistics and pin-point regions of active star formation, while the multi-wavelength photometry reveal the relative distributions of pre-stellar cores, protostars and pre-main-sequence objects. Benefits are also wrought when \htwo\ data are combined with optical and radio data, in outflows from low and high mass YSOs, but also toward post-AGB stars and PN, where rapidly changing physical conditions require combined observations in a variety of excitation tracers. The GLIMPSE-North data are already available and accessible through the IPAC archive (and has been cross-matched to the GPS, available at the WFCAM science archive); IPHAS data are also publicly available and cover half of the Northern Galactic Plane; the AKARI all sky survey faint point source catalogue will also be available in the near future. The MSX point-source catalogue may also be used for the few massive sources that saturate in IRAC data, along with future longer wavelength data from JCMT-JPS. The importance of serendipitous discoveries of interesting emission-line objects of all types should also not be ignored. The combination of broad- and narrow-band data will aid in the selection of targets for study with instruments where full-scale mapping of the Galactic plane is impractical (e.g. the sub-mm heterodyne array receiver, HARP, the SCUBA-2 Polarimeter on the JCMT, and of course future high-resolution facilities such as ALMA and the James Webb Space telescope). Finally, it is worth noting that massive star formation is at the heart of the science case for many Galactic Plane surveys, e.g. the JCMT Galactic Plane Survey (JPS), the Spitzer-GLIMPSE survey, and the Herschel Hi-GAL mid-IR survey. These will reveal dense molecular cores and mid-IR bright YSOs, effectively identifying all of the massive pre-main-sequence stars, protostars and massive pre-stellar cores. UWISH2 will in turn establish how many of these are dynamically active. Although extinction will be high in the coldest and most massive objects, outflows rapidly break out of these environments and are usually bright in \htwo\ line emission. Identifying outflows and tracing them back to sub-mm sources will certainly be an important way of identifying the location of protostars in these mid-IR and sub-mm survey data. Indeed, our initial analysis of 24 square degrees of early UWISH2 data, which has revealed numerous outflows (as well as new PN, clusters and variable stars), demonstrates that this is certainly the case.

1012.0782

1012

1012.0217_arXiv.txt

{We study the capabilities of the {\it Fermi}--LAT instrument for identifying particle Dark Matter properties as mass, annihilation cross section and annihilation channels, with gamma-ray observations from the Galactic Center. For the potential Dark Matter signal, besides the prompt gamma-ray flux produced in Dark Matter annihilations, we also take into account the flux produced by inverse Compton scattering of the electrons and positrons generated in Dark Matter annihilations off the interstellar photon background. We show that the addition of this contribution is crucial in the case of annihilations into $e^+e^-$ and $\mu^+\mu^-$ pairs. In addition to the diffuse galactic and extragalactic background, we also consider the full catalog of high-energy gamma-ray point sources detected by {\it Fermi}. The impact of the degeneracies between the different Dark Matter annihilation channels has been studied. We find that for Dark Matter masses below $\sim 200$ GeV and for typical thermal annihilation cross sections, it will be possible to obtain stringent bounds on the Dark Matter properties.} \FullConference{Identification of Dark Matter 2010\\ July 26 - 30 2010\\ University of Montpellier 2, Montpellier, France} \begin{document}

If Dark Matter (DM) is detected and identified, the measurement of its properties like mass, annihilation cross-section and annihilation channels plays a central role in the determination of the particle nature of the DM. It will allow us to constrain models of particle physics beyond the Standard Model, for instance supersymmetry and universal extra dimensions. Furthermore a convincing DM discovery may require consistent signals in multiple experiments in multiple channels (direct, indirect, collider). We discuss the capabilities of the {\it Fermi}--LAT instrument for identifying particle DM properties with gamma-ray observations from the Galactic Center (GC). The differential intensity of the photon signal from a given observational region in the galactic halo from the annihilation of DM particles has different possible origins: internal bremsstrahlung and secondary photons (prompt) as well as Inverse Compton Scattering (ICS). External bremsstrahlung and synchrotron emission also contribute to the photon flux; however, for the energies of interest here and for typical DM masses, both bremsstrahlung and synchrotron emission are expected to be subdominant with respect to ICS. For the sake of simplicity we will neglect these sources in what follows. The differential flux of prompt gamma-rays from DM annihilations and coming from a direction within a solid angle $\Delta\Omega$ is given by \begin{equation} \left(\frac{d\Phi_{\gamma}}{dE_\gamma}\right)_{{\rm prompt}} (E_{\gamma},\, \Delta\Omega) = \frac{\langle\sigma v\rangle}{2\,m_\chi^2}\sum_i\frac{dN_{\gamma}^i}{dE_{\gamma}}\, \textrm{BR}_i \, \frac{1}{4\,\pi} \, \int_{\Delta\Omega}d\Omega \, \int_\textrm{los}\rho\big(r(s,\,\Omega)\big)^2 \, ds\,, \label{Eq:promptflux} \end{equation} where $\langle\sigma v\rangle$ is the total thermally averaged annihilation cross section, $m_\chi$ the mass of the DM particle, $\textrm{BR}_i$ the annihilation fraction into channel $i$, $dN_\gamma^i/dE_\gamma$ the differential gamma-ray yield of standard model particles into photons of energy $E_\gamma$, $\rho(r)$ the DM density profile and $r$ the distance from the GC. Here we will focus on the NFW halo profile \cite{Navarro:1995iw}; the dependence on the DM halo profile has been studied in reference \cite{Bernal:2010ip}. An abundant population of energetic electrons and positrons produced in DM annihilations either directly or indirectly from the hadronization, fragmentation, and subsequent decay of the SM particles in the final states, gives rise to secondary photons at various wavelengths via ICS off the diffuse radiation fields in the galaxy. We approximate this photon background as a superposition of three black-body spectra consisting of the CMB, the optical starlight and the infrared radiation due to rescattering of starlight by dust \cite{ics}. The differential flux of high energy photons produced by the ICS processes is given by \cite{Blumenthal:1970gc} \begin{equation} \left(\frac{d\Phi_{\gamma}}{dE_\gamma}\right)_{{\rm ICS}} (E_{\gamma},\, \Delta\Omega) = \frac{1}{E_{\gamma}} \, \frac{1}{4\pi} \, \int_{\Delta\Omega}d\Omega \, \int_\textrm{los}ds \, \int_{m_e}^{m_\chi}dE\,\mathcal{P}(E_{\gamma},\,E) \, \frac{dn_e}{dE}\big(E,\,r,\,z\big) ~, \label{Eq:ICSflux} \end{equation} where $\mathcal{P}(E_{\gamma},\,E)$ is the differential power emitted into scattered photons of energy $E_{\gamma}$ by an electron with energy $E$. The minimal and maximal energies of the electrons are determined by the electron mass $m_e$ and the DM particle mass. The quantity $dn_e/dE$ is the electron plus positron spectrum after propagation in the Galaxy, which will differ from the energy spectrum produced at the source. We determine the propagated spectrum by solving the diffusion-loss equation that describes the evolution of the energy distribution for electrons and positrons assuming steady state~\cite{propaga}. Regarding the propagation parameters (like diffusion coefficient, energy losses and thickness of the diffusion zone), we take their values from the commonly used MED model~\cite{propaga}. Again, the dependence on the propagation model has been studied in reference \cite{Bernal:2010ip}. There are three main components contributing to the high-energy gamma-ray background: the diffuse galactic emission has been estimated by taking the conventional model of the GALPROP code~\cite{galprop}. On the other hand, another source of background particularly important when looking at the GC is that of resolved point sources. We consider all the point sources detected by the first 11~months of {\it Fermi}--LAT~\cite{pointsources} lying in the region of interest. Finally, for the isotropic extragalactic gamma-ray background we used the recent measurements by the {\it Fermi}--LAT collaboration~\cite{Abdo:2010nz}. For a $10^\circ \times 10^\circ$ region around the GC, the diffuse galactic emission dominates below $\sim 20$~GeV. Above that value, the emission coming from point sources is the most important. The isotropic extragalactic gamma-ray background is at the percent level. The Large Area Telescope ({\it Fermi}--LAT) is the primary instrument on board of the {\it Fermi Gamma-ray Space Telescope}. It performs an all-sky survey, covering a large energy range for gamma-rays, with an effective area $\simeq 8000$~cm$^2$ and a field of view of $2.4$~sr. In the following analysis, we consider a 5-year mission run, and an energy range from 1~GeV extending up to 300~GeV. We divide this energy interval into 20 evenly spaced logarithmic bins. In order to maximize the signal-to-noise ratio, it has been pointed out that for a NFW profile the best strategy is to focus on a region around the GC of $\sim10^\circ\times 10^\circ$~\cite{maxsb}. Hence, this is our choice.

In this work we have studied the abilities of the {\it Fermi}--LAT instrument to constrain Dark Matter properties by using the current and future observations of gamma-rays from the Galactic Center produced by DM annihilations. Unlike previous works, we also take into account the contribution to the gamma-ray spectrum from ICS of electrons and positrons produced in DM annihilations off the ambient photon background. We show that the inclusion of the ICS contribution for hadronic channels and for the $\tau^+\tau^-$ channel does not give rise to important differences in the reconstruction process. This is not the case if DM annihilates into the $\mu^+\mu^-$ and the $e^+e^-$ channel. In this latter case, adding the ICS contribution to the prompt gamma-ray spectrum turns out to be crucial in order not to obtain completely wrong results. On the other hand, we found that for Dark Matter masses below $\sim 200$ GeV and for typical thermal annihilation cross sections, it will be possible to obtain stringent bounds on the Dark Matter properties such as its mass, annihilation cross section and annihilation channels.

1012.0217

1012

1012.1891.txt

%This paper presents a statistical analysis of the mid-infrared (MIR) spectra of 248 luminous infrared (IR) galaxies (LIRGs) which comprise the Great Observatories All-sky LIRG Survey (GOALS) observed with the Infrared Spectrograph (IRS) on-board the Spitzer Space Telescope in the rest-frame wavelength range between 5 and 38 $\mu$m. The GOALS sample enables a direct measurement of the relative contributions of star-formation and active galactic nuclei (AGN) to the total IR emission from a large, statistically complete sample of LIRGs in the local Universe. %The AGN contribution to the MIR emission is estimated by employing several diagnostics based on the properties of the [NeV], [OIV] and [NeII] fine structure gas emission lines, the 6.2 $\mu$m PAH equivalent width (EQW) as well as the shape of the MIR continuum. The [NeV] line, which indicates the presence of an AGN, is detected in 18\% of all LIRGs. The 6.2 $\mu$m PAH EQW, [NeV]/L$_{\rm{IR}}$, [NeV]/[NeII] and [OIV]/[NeII] ratios, and the ratios of 6.2 $\mu$m PAH flux to the integrated continuum flux between 5.3 and 5.8 $\mu$m suggest that in 10\% of the sources the AGN contributes more than 50\% of the total IR luminosity. When summing up the total IR luminosity contributed by AGN in all the LIRGs in our sample, we find that AGN contribute approximately 12\% to the total energy emitted by LIRGs in the local universe. %We find that average spectrum of sources with an AGN looks qualitatively similar to the average spectrum of sources without an AGN. However the average spectrum of sources in which the AGN contributes at least a third of the MIR emission have lower PAH emission and flatter MIR continua than sources without an AGN. We also find that AGN dominated LIRGs have higher total and nuclear (i.e. within the few central kilo-parsecs corresponding to the projected size of the IRS slit) IR luminosities, warmer MIR colors and are found in interacting systems more often than starburst dominated LIRGs. However there are no obvious linear correlations between these properties, implying that none of these properties alone can be used to quantify exactly the activity and evolution of an individual LIRG. %A study of the IRAC colors of LIRGs confirms that methods of finding AGN on the basis of their MIR colors are effective at choosing AGN but 40\% to 50\% of AGN dominated LIRGs are not selected as such with these methods. We present a statistical analysis of the mid-infrared (MIR) spectra of 248 luminous infrared (IR) galaxies (LIRGs) which comprise the Great Observatories All-sky LIRG Survey (GOALS) observed with the Infrared Spectrograph (IRS) on-board the Spitzer Space Telescope. The GOALS sample enables a direct measurement of the relative contributions of star-formation and active galactic nuclei (AGN) to the total IR emission from a large sample of local LIRGs. The AGN contribution to the MIR emission (f$_{\rm{AGN}}$) is estimated by employing several diagnostics based on the properties of the [NeV], [OIV] and [NeII] fine structure gas emission lines, the 6.2 $\mu$m PAH and the shape of the MIR continuum. We find that 18\% of all LIRGs contain an AGN and that in 10\% of all sources the AGN contributes more than 50\% of the total IR luminosity. Summing up the total IR luminosity contributed by AGN in all our sources suggests that AGN supply $\sim12\%$ of the total energy emitted by LIRGs. The average spectrum of sources with an AGN looks similar to the average spectrum of sources without an AGN, but it has lower PAH emission and a flatter MIR continuum. AGN dominated LIRGs have higher IR luminosities, warmer MIR colors and are found in interacting systems more often than pure starbursts LIRGs. However we find no linear correlations between these properties and (f$_{\rm{AGN}}$). We used the IRAC colors of LIRGs to confirm that finding AGN on the basis of their MIR colors may miss $\sim 40$\% of AGN dominated (U)LIRGs.

The Infrared (IR) Astronomical Satellite (IRAS) provided the first unbiased survey of the sky at mid-infrared (MIR) and far-infrared (FIR) wavelengths, giving us a comprehensive census of the IR emission properties of galaxies in the local Universe. IR number counts have been used to trace the importance of IR emission as a function of redshift, to explore star-formation and galaxy evolution \citep[]{flores99, gispert00, franceschini01, chary01, chary04, elbaz02, lagache03, marleau04, lefloc05, caputi06, magnelli09}. To understand the origin of the observed IR emission, MIR diagnostic tools based on ISO \citep[for an exhaustive review see:][]{gec00} have been developed to study the roles played by star formation, AGN, and shocks (interaction-driven and wind-driven) in producing the observed IR emission. These diagnostics permit a direct mapping between IR number counts as a function of redshift/luminosity/galaxy type and the evolution of accretion and star-formation. The most basic diagnostics employed to estimate the AGN contribution to the MIR emission in individual galaxies are the ratios of high to low ionization fine-structure emission lines: [NeV] 14.3 $\mu$m/[NeII] 12.8 $\mu$m and [OIV] 25.9 $\mu$m/[NeII] 12.8 $\mu$m. The [NeV] 14.3 $\mu$m and [OIV] 25.9 $\mu$m lines trace high ionization gas. These methods have been used by e.g. \citet{lutz99, genzel98, verma03, armus06, sturm02, spoon07, armus07, farrah07}. The ionization potential of [NeV] is 96 eV. This is too high to be produced by OB stars, therefore its detection in the integrated spectrum of a galaxy usually indicates the presence of an AGN. This is not true for the [OIV] line because it takes only 55 eV to ionize O$^{++}$. Empirically it has been shown that emission line ratios of [NeV]/[Ne II] $\geq$ 0.75 and [OIV]/[NeII] $\geq 1.75$ indicate that more than 50\% of the nuclear MIR emission is produced by an AGN \citep[e.g.][and ref. within]{armus07}. Additional diagnostics of the relative contribution of starbursts (SB) and AGN to the MIR luminosity are based on the dust properties, in particular PAH emission lines and the continuum emission arising from dust being heated by a SB or an AGN. Mid-infrared continuum emission in galaxies arises from a combination of ionized interstellar gas, evolved stellar population, non-thermal emission from radio sources, very small grains and PAHs. Empirically their respective contributions can be roughly distinguished from the shape of the SED \citep{laurent00}. Both theoretical models \citep[e.g.][]{pier92, nenkova02, granato07, levenson07} and observations \citep[e.g.][] {alonso01} show that the AGN radiation field can heat grains such that the dust continuum emission becomes prominent between 3 and 6 $\mu$m. Galaxies with an AGN tend to have low 6.2 $\mu$m PAH equivalent width (EQW) \citep[e.g][]{genzel98, lutz99, rigop99, tran01, sturm00, desai07, wu09} due to the presence of a significant hot dust continuum and also because the hard AGN photons may destroy the PAH molecules. %In particular the Infrared Spectrograph \citep[IRS][]{houck04} on Spitzer has led to significant progress in our understanding of low luminosity starbursts \citep[SBs, e.g.][]{brandl06, dale06, smith07} and AGN \citep[e.g.][]{weedman05, gorjian07, hao07, wu09}, ULIRGs \citep[e.g.][]{armus04, armus06, armus07, desai07, farrah07, iman08, lahuis07, spoon04, spoon07, veilleux09a} and small samples of local LIRGs \citep[e.g][]{pere10, tanio10, tanio08, inami 2010, evans08} with the study of large sets of those objects. Spitzer IRS observations of normal nearby optically classified galaxies confirmed that the MIR diagnostics described above, diagnostics based on a combination of high to low ionization line ratios and PAH strengths are indeed very effective at determining the AGN versus starburst contribution to the IR in nearby galaxies. These studies such as the Spitzer Infrared Nearby Galaxy Survey (SINGS) were based on galaxies from a wide range of environments, Hubble Types and dust contents \citep{dale06}. %, but that it is difficult to distinguish between pure Seyferts and LINERs using these diagnostics. Most recently \citet{goulding09} looked at an optically selected nearby ($D<15 $ Mpc) sample of galaxies with IR luminosity between $3 \times 10^{9} ~\rm{and}~ 2\times 10^{11} L_{\odot}$ and found that 27\% have AGN based on [NeV] detections. Excluding the nearby LIRG NGC1068, the [NeV]/[NeII] ratios and the PAH EQWs of these sources suggest that only ~5\% of sources are AGN dominated. IRS studies of 24 nearby starbursts with luminosities between 10$^{9.75}$ and 10$^{11.6}$ at an average distance of 33 Mpc %and average IR luminosity $5\times 10^{10} L_{\odot}$ provided high and low resolution spectral templates for pure starbursts %by excluding composite (AGN + SB) objects where the evidence for the presence of an AGN came from literature and from the detection of the [NeV] line \citep{bers09,brandl06}. \citet{brandl06} also found the spectral continuum slope longward of $15 ~\mu$m can be used to discriminate between SB and AGN powered sources. These authors also found that in pure starbursts the PAH EQWs are independent of $L_{IR}$ but that the luminosity of the PAH feature scales with $L_{IR}$, in particular the 6.2 $\mu$m feature can be used to approximate the total IR of the starburst. %In addition to the work determining the zero point for galaxies where the starburst emits all the observed IR emission % %QSOs Weedman et al. 2005 apj 633, 706 %Sample: 8 classic but diverse AGN %Method: Continuum slopes, high ionization lines, PAH EQWs

This paper presents a statistical analysis of 248 LIRG spectra in the rest-frame wavelength range between 5 and 38 $\mu$m. Several diagnostics effective at isolating the Active Galactic Nuclei (AGN) contribution to the Mid-infrared (MIR) emission using [NeV], [OIV] and [NeII] gas lines, the 6.2 $\mu$m PAH EQW and the shape of the MIR continuum are compared. In summary: \begin{enumerate} \item The high ionization emission lines of [NeV] 14.322 $\mu$m and [OIV] 25.890 $\mu$m are detected in 18\% and 53\% of all LIRG nuclei respectively.% while the 6.2 $\mu$m PAH and [NeII] 12.8 $\mu$m features are detected in 98\% and 100\% of the sample. Since the [NeV] line does not arise from gas heated by hot stars, its detections suggest the presence of an AGN in at least 18\% of LIRG nuclei. \item Diagnostics using the [NeV]/[NeII], [OIV]/[NeII] line flux ratios, the 6.2 $\mu$m PAH EQW and the MIR continuum shape suggest that in 10\% of local LIRGs the AGN dominates the bolometric luminosity. The vast majority of LIRGs are SB dominated. The fraction of local LIRG IR emission coming from an AGN as estimated in the mid infrared is approximately two times larger than that seen in normal galaxies ($\sim$5\%) and, about three to four times lower than that seen in ULIRGs alone. \item Summing the bolometric luminosity contributed by each AGN in the sample and dividing by the total IR luminosity of all the LIRGs suggests that AGN are responsible for $\sim$12\% of the total bolometric luminosity of local LIRGs. \item In LIRGs there are no strong correlations between the fraction of IR luminosity from an AGN and the total or nuclear 24 $\mu$m luminosity, the 24 to 60 $\mu$m flux ratios or the interaction stage of the system. However AGN dominated LIRGs tend to be more luminous at 24 $\mu$m and to have warmer IR colors than starburst dominated LIRGs. \item By separating the GOALS sources according to merger stage it is found that there is a significant increase in the fraction of AGN dominated sources among those galaxies in the latest stages of interaction. This trend is driven by the ULIRGs in the sample, since these objects tend to be late stage mergers and have larger AGN fractions than the LIRGs . This is consistent with findings of previous authors using optical diagnostics for LIRGs, MIR studies of ULIRGs and PG QSOs and with models which predict that mergers of gas-rich spirals fuel both star-formation and accretion onto a super-massive black hole. \item An investigation of the IRAC colors (i.e. [3.6 $\mu$m]-[4.5 $\mu$m] versus [5.8$\mu$m]-[8 $\mu$m]), as introduced in \citet{stern05}, of LIRGs indicates that only 50\% of objects with a significant AGN contribution to the MIR emission fall within the range typically associated with AGN. The \citet{lacy04} AGN IRAC color criteria select a slightly higher fraction of the AGN dominated LIRGs (64\%), at the expense of also including 11 LIRGs that appear SB dominated from their IRS spectra. \end{enumerate} % IR number counts have been used by several groups of researchers to determine the star-formation history of our universe. The measurements provided in this paper are important in this context because they provide an estimate of the fraction of local LIRG IR emission coming from an AGN and thus not from star-formation. If the physical properties of LIRGs do not strongly evolve between redshifts 0 and 1, then the statistics derived here can be used to ascertain the properties IR luminous systems at redshfits where their numbers as well as the contribution to the total energy production in the universe is substantial. In particular, these results can be used as a correction when determining the star formation rate from IR number counts as a function of redshift. The measurements we present in this paper provide an estimate of the fraction of IR emission in LIRGs coming from AGN. Our results provide an important local benchmark against which to compare high-redshift samples of LIRGs, especially at epochs where the contribution of LIRGs to the IR background (e.g. at z$\sim$1 - see Magnelli et al. 2009) becomes substantial. These diagnostics probe the nuclear source of IR emission in those LIRGs. A full understanding of the processes leading to the generation of LIRG activity requires careful analysis of the mass, temperature and kinematics of the gas fueling and being heated by the star-formation and AGN activity in LIRGs. Future papers (Petric et al. in prep) will discuss the observations of warm and cold molecular gas in the GOALS sample and relate these to the energy sources and evolutionary state of the LIRGs. % = = = = = %TABLES % = = = = = \begin{deluxetable}{lcccccc} %\rotate \tableheadfrac{0.0} \tablecolumns{5} \tabletypesize{\scriptsize} \tablewidth{0pt} \tablecaption{[NeV] 14.3 $\mu$m detections in the GOALS sample} \tablehead{ \colhead{Name} &\colhead{RA J2000} &\colhead{Dec J2000} &\colhead{Flux\tablenotemark{*} [NeV]} &\colhead{Error } &\colhead{${[NeV]}\over{[NeII]}$} &\colhead{${L_{[NeV]}}\over{L_{IR}}$}\\ \colhead{} &\colhead{[deg]} &\colhead{[deg]} &\colhead{$1.0\times 10^{-18}$} &\colhead{$1.0\times 10^{-18}$ } &\colhead{$1.0\times 10^{-2}$} &\colhead{$1.0\times 10^{-4}$}\\ } \startdata NGC 0232 & 10.7201 & -23.541 & 38.4 & 5.4 & 31.20 & 1.03 \\ Mrk 1034 & 35.8416 & 32.1971 & 22.6 & 3.3 & 6.31 & 0.33 \\ NGC 1068$^{a}$ & 40.6696 &-0.0133 & 9150.0 & 180.0 & 209.00 & 1.80 \\ UGC 02608 & 48.7561 &42.0357 & 318.00 & 97.1 & 52.5 & 3.67 \\ NGC 1365$^{b}$ & 53.4015 & -36.1404 & 216.0 & 10.5 & 13.40 & 0.30 \\ ESO 420-G013 & 63.4571 & -32.0070 & 113.0 & 7.9 & 9.60 & 0.76 \\ UGC 03094 & 68.8910 & 19.1717& 15.3 & 3.0 & 4.79 & 0.20 \\ CGCG 468-002NED01 &77.0821 & 17.3633& 21.3 & 1.8 & 28.10 & 0.54 \\ IRAS F05081+7936 & 79.1933 &79.6703 & 22.6 & 3.7 & 4.40 & 0.46 \\ IRAS F05189-2524$^{c}$ & 80.2559 & -25.3626 & 151.0 & 11.8 & 85.00 & 1.06 \\ IRAS 05083+244 & 77.8578 &24.7551 & 11.3 & 3.9 & 1.28 & 0.18\\ IRAS F07027-6011 & 105.8506 &-60.2561 & 26.9 & 2.1 & 19.30 & 0.44 \\ IRAS 16164-0746 & 244.7991 &-7.9008& 11.9 & 2.9 & 2.40 &0.18\\ NGC 2623 & 129.6003 & 25.7547 & 26.2 & 4.0 & 4.82 & 0.13 \\ 2MASX J09133644-1019296 &138.4021 & -10.3249 & 12.5 & 1.6 & 9.84 & 0.23 \\ 2MASX J09133888-1019196 &138.4120 & -10.3221 & 13.8 & 1.6 & 10.80 & 1.17 \\ UGC 5101$^{d,f}$ & 143.9652 & 61.353 & 28.6 & 1.9 & 7.08 & 0.22\\ ESO 267-G030 & 183.5534 & -47.2285 & 19.4 & 3.7 & 4.41 & 0.28 \\ NGC 4922NED02 & 195.3553 & 29.3138 & 30.8 & 6.4 & 8.10 & 0.41 \\ UGC 08387 & 200.1473 & 34.1395 & 14.8 & 4.4 & 1.33 & 0.09 \\ MCG -03-34-064 & 200.6019 & -16.7284 & 547.0 & 19.7 & 112.00 & 6.55 \\ NGC 5135$^{e}$ & 201.4332 &-29.8333 & 106.0 & 6.3 & 11.70 & 0.45 \\ Mrk 266$^{b}$ & 204.5737 & 48.2761 & 21.5 & 2.9 & 14.05 & 0.28 \\ NGC 5256 & 204.5719 & 48.2756 & 79.6 & 12.8 & 13.70 & 1.42 \\ Mrk 273$^{b}$ & 206.1755 & 55.8870 & 101.0 & 3.3 & 23.70 & 0.66 \\ NGC 5734 & 221.2959 & -20.9135 & 9.9 & 3.3 & 2.90 & 0.10 \\ VV340a & 224.2529 & 24.6183 & 13.3 & 2.1 & 4.74 & 0.16 \\ NGC 5990 & 236.5684 & 2.4154 & 12.0 & 4.1 & 1.94 & 0.08 \\ NGC 6156 & 248.7190 & -60.6189 & 11.9 & 4.2 & 4.11 & 0.06 \\ NGC 6240$^{f}$ & 253.2454 & 2.4009 & 20.7 & 7.2 & 1.17 & 0.07 \\ CGCG 141-034 & 269.2360 & 24.0172 & 20.3 & 1.8 & 5.40 & 0.28 \\ CGCG 142-034B & 274.1410 & 22.1108 & 10.1 & 2.7 & 8.30 & 0.15 \\ ESO 339-G011 & 299.4067 & -37.9357 & 245.0 & 4.1 & 55.30 & 3.06 \\ MCG +04-48-002 & 307.1461 & 25.7334 & 29.0 & 5.6 & 5.34 & 0.28 \\ NGC 6926 & 308.2755 & -2.0275 & 13.4 & 2.4 & 19.30 & 0.13 \\ NGC 7130 & 327.0813 & -34.9517& 58.6 & 4.6 & 8.79 & 0.27 \\ NGC 7469 & 345.8151 & 8.8740 & 119.0 & 15.4 & 5.75 & 0.38 \\ NGC 7679 & 352.1943 & 3.5115 & 32.9 & 4.2 & 5.2 & 0.39 \\ CGCG 453-062 & 346.2355 & 19.5522 & 22.7 & 6.9 & 7.95 & 0.29 \\ ESO 148-IG002 & 348.9459 & -59.0547 & 20.4 & 3.9 & 6.59 & 0.19 \\ IC 5298 & 349.0029 & 25.5567& 105.0 & 4.9 & 30.50 & 0.79 \\ NGC 7592 & 349.5946 & -4.4162 & 14.9 & 1.5 & 3.60 & 0.21 \\ NGC 7674 & 351.9863 & 8.7790 & 186.0 & 7.8 & 97.30 & 2.16 \\ % % % NGC0235A & 38.40 & 5.41 & 31.20 & 1.03 \\ % Mrk1034 & 22.6 & 3.3 & 6.31 & 0.33 \\ % % %NGC0876 & 2.59 & 6.0 & 1.95 & 0.06 \\ % NGC1068$^{a}$ & 9150.00 & 180.0 & 209.00 & 1.80 \\ % UGC02608 & 318.00 & 97.1 & 52.50 & 3.67 \\ % NGC1365$^{b}$ & 216.00 & 10.53 & 13.40 & 0.30 \\ % ESOO420-G013 & 113.00 & 7.93 & 9.60 & 0.76 \\ % UGC03094 & 15.30 & 2.98 & 4.79 & 0.20 \\ % CGCG468-002NED01 & 21.30 & 1.81 & 28.10 & 0.54 \\ % IRASF05081+7936 & 22.60 & 3.71 & 4.40 & 0.46 \\ % IRASF05189-2524$^{c}$ & 151.00 & 11.75 & 85.00 & 1.06 \\ % IRAS05083+244 & 11.30 & 3.93 & 1.28 & 0.18\\ % IRASF07027-6011 & 26.90 & 2.07 & 19.30 & 0.44 \\ % iras16164-0746 & 11.87 & 2.87 & 2.40 &0.18\\ % NGC2623 & 26.20 & 4.03 & 4.82 & 0.13 \\ % 2MASXJ09133644-1019296 & 12.50 & 1.55 & 9.84 & 0.23 \\ % 2MASXJ09133888-1019196 & 13.80 & 1.55 & 10.80 & 1.17 \\ % UGC05101$^{d,e}$ & 28.6 & 1.91 & 7.08 & 0.22 \\ % %ESO264-G036 & 3.13 & 0.5 & 1.93 & 0.04 \\ % ESO267-G030 & 19.40 & 3.7 & 4.41 & 0.28 \\ % NGC4922NED02 & 30.80 & 6.37 & 8.10 & 0.41 \\ % UGC08387 & 14.80 & 4.42 & 1.33 & 0.09 \\ % MCG-03-34-064 & 547.00 & 19.73 & 112.00 & 6.55 \\ % NGC5135$^{f}$ & 106.00 & 6.3 & 11.70 & 0.45 \\ % Mrk266$^{b}$ & 21.5 & 2.86 & 14.05 &0.28\\ % NGC5256 & 79.60 & 12.8 & 13.70 & 1.42 \\ % Mrk273$^{b}$ & 101.00 & 3.33 & 23.70 & 0.66 \\ % NGC5743 & 9.90 & 3.27 & 2.90 & 0.10 \\ % VV340a & 13.30 & 2.11 & 4.74 & 0.16 \\ % NGC5990 & 11.99 & 4.07 & 1.94 & 0.08 \\ % %IZw107 & 16.77 & 2.93 & 5.59 &0.19\\ % NGC6156 & 11.90 & 4.16 & 4.11 & 0.06 \\ % NGC6240$^{g}$ & 20.66 & 7.22 & 1.17 & 0.07 \\ % CGCG141-034 & 20.30 & 1.75 & 5.40 & 0.28 \\ % CGCG142-034B & 10.10 & 2.71 & 8.30 & 0.15 \\ % ESO339-G011 & 245.00 & 4.07 & 55.30 & 3.06 \\ % MCG+04-48-002 & 29.00 & 5.57 & 5.34 & 0.28 \\ % NGC6926 & 13.40 & 2.41 & 19.30 & 0.13 \\ % NGC7130 & 58.60 & 4.57 & 8.79 & 0.27 \\ % NGC7469 & 119.00 & 15.4 & 5.75 & 0.38 \\ % NGC7679 & 32.87 & 4.2 & 5.2 & 0.39\\ % CGCG453-062 & 22.70 & 6.93 & 7.95 & 0.29 \\ % ESO148-IG002 & 20.40 & 3.9 & 6.59 & 0.19 \\ % IC5298 & 105.00 & 4.90 & 30.50 & 0.79 \\ % NGC7592& 14.93 & 1.53 &3.60& 0.21\\ % NGC7674 & 186.00 & 7.8 & 97.30 & 2.16 \\ \enddata \tablenotetext{*}{Notes: All fluxes are given in W m$^{-2}$. The fluxes presented here agree with previously published values within the errors, in all cases with the exception of fluxes for Mrk 266, NGC5135. } \tablenotetext{a}{Howell et al. 2007}%\citep{howell07}} \tablenotetext{b}{Dudik et al. 2007}%\citep{dudik07}} \tablenotetext{c}{Armus et al. (2007)} \tablenotetext{d}{Farrah et al. (2007)} %\tablenotetext{e}{Farrah et al. (2007)} \tablenotetext{e}{Gorjian et al. (2007)}%\citep{gorjian07}} \tablenotetext{f}{Armus et al. 2006} \end{deluxetable} % = = = = = %FIGURES % = = = = = % \begin{figure}[htb] % \begin{center} %\end{center} % \includegraphics[width=\linewidth]{Ice_ex.jpg} % \caption{IRAC color ([5.8 $\mu$m ] - [8 $\mu$m]) color ([3.6 $\mu$m ] - [4.5 $\mu$m]) plot of 224 LIRG nuclei. The IRAC total fluxes shown here are from apertures matched to the MIPS 24 $\mu$m apertures % of typically 1', but range between 0.5' to 1.5', (Mazzarella et al. 2010 in prep.). The solid lines show the color cuts which Stern et al. (2005) use to separate active galaxies from Galactic stars and normal galaxies. The solid black, red and blue dots represent points with 6.2 $\mu$m PAH EQW greater than 0.27 $\mu$m, between 0.14 and 0.27 $\mu$m and smaller than 0.14 $\mu$m respectively. The empty circles represent nuclei without MIR spectra.} % \label{ice_ex} % \end{center} % \end{figure} \begin{figure}[htb] \begin{center} \includegraphics[width=\linewidth]{petric_fig1.eps} %Nov01_ne5_diag0.eps l 2821} \caption{Mid-infrared [NeV]/[NeII] versus 6.2 $\mu$m PAH EQW excitation diagram. The \textit{red triangles} are [NeV] detections while the \textit{black arrows} are upper limits. In all cases 1$\sigma$ error bars are shown.The black circles indicate ULIRGs where the filled symbols show detections and empty symbols upper limits. The blue empty stars mark upper limits for star-burst galaxies from \citet{bers09}. The \textit{solid black lines} indicate the fractional AGN and starburst contribution to the MIR luminosity from the [NeV]/[NeII] (vertical) and 6.2 $\mu$m PAH EQW (horizontal) assuming a simple linear mixing model. In each case, the 100$\%$, 50$\%$, 25$\%$, and 10$\%$ levels are marked. The 100$\%$ level is set by the average detected values for the [NeV]/NeII] and 6.2 $\mu$m EQW among AGN and starbursts respectively, as discussed in \citet{armus07}. The \textit{blue line} traces where the summed SB and AGN contribution equals 100$\%$. For most LIRGs the [NeV]/[NeII] ratio suggests that the AGN contribution to the nuclear MIR luminosity is below 10$\%$. } \label{Nov01_ne5_diag0} \end{center} \end{figure} \begin{figure}[htb] \begin{center} \includegraphics[width=\linewidth]{petric_fig2.eps}% Nov01_O4ne2_basic_diag_main_2.eps} \caption{Mid-infrared [OIV]/[NeII] versus 6.2 $\mu$m PAH EQW excitation diagram. Symbols have the same definition as in Figure 1 except that red triangles indicate [OIV] detections} \label{O4ne2_basic_diag} \end{center} \end{figure} \begin{figure}[htb] \begin{center} \includegraphics[width=\linewidth]{petric_fig3.eps}% Laurent_plot_stage.eps} \caption{Mid-infrared diagnostic diagram, first used by Laurent et al. (2000) and later modified for Spitzer IRS by Armus et al. (2007), comparing the integrated continuum flux from 14-15 $\mu$m, the integrated continuum flux from 5.3-5.5 and the 6.2 $\mu$m PAH flux. The three vertices, labeled as AGN, HII and PDR, represent the positions of 3C273 from \citet{weedman05}, M17 and NGC 7023 from \citet{peeters04}. These vertices were chosen to facilitate comparison with ULIRGs as presented in \citet{armus07} and \citet{brandl06}. The red lines from left to right indicate a 90$\%, ~ 75\%$ and 50$\%$ fractional AGN contribution to the nuclear MIR luminosity respectively. GOALs sources are shown as open circles.} \label{Laurent-plot} \end{center} \end{figure} % \begin{figure}[htb] % \begin{center} % \includegraphics[width=\linewidth]{Ave_spec.eps} % \caption{Top Panel: Average spectra of galaxies whose [NeV]/[NeII] and [OIV]/[NeII] ratios suggest they contain AGNs which contribute more than 10\% to the MIR emission. The spectra were normalized by the flux at 24 $\mu$m and were weighted by the signal to noise ratio (SNR) at 24 $\mu$m. The error-bars give the 1$\sigma$ error on the average while the shaded region shows the intrinsic weighted dispersion in the spectra that were combined to determine the average. Middle panel: Average spectra of galaxies where the same diagnostics indicate the the MIR emission originates in star-formation regions. Bottom panel: comparison of the two average spectra presented in the top two panels. As expected the PAH emission is significantly stronger in the second average (i.e. PAH EQWs a factor of 2-3 larger). The 6.2 $\mu$m PAH EQW in sources where the MIR comes from star-formation is twice that of galaxies with an AGN contribution to the MIR emission larger than 10\%. The continuum slopes as estimated from the fluxes at 30 $\mu$m, 15 $\mu$m and 5.5 $\mu$m are slightly rising with wavelength for the former but are flat for the latter. \citet{veilleux09a} find that the 30 $\mu$m to 15 $\mu$m flux ratio is a powerful continuum diagnostics of AGN activity among ULIRGs and QSOs. This figure suggests that the same diagnostic can be used to roughly assess the average AGN contribution to the IR emission in LIRGs.} % \label{Avg_spectra} % \end{center} % \end{figure} \begin{figure}[htb] \begin{center} \includegraphics[width=\linewidth]{petric_fig4.eps}%Ave_specs.eps} \caption{Average low-resolution spectra of 4 groups of galaxies. The spectra were normalized by the flux at 24 $\mu$m and were weighted by the signal to noise ratio (SNR) at 24 $\mu$m. The error-bars give the 1$\sigma$ error on the average while the shaded region shows the intrinsic weighted dispersion in the spectra that were combined to determine the average. From top to bottom we show averages of the following groups of objects: Group (1) contains sources without detectable [NeV] and with [OIV]/[NeII] flux ratios $\leq 0.35$, that is sources whose MIR luminosity is dominated by star-formation. Group (2) contains sources with [NeV] detections. Group (3) contains sources with [NeV]/[NeII] $\geq$ 0.14 or [OIV]/[NeII] flux ratios $\geq$ 0.5 suggesting an AGN contribution to the MIR greater than $\sim$ 33\%. Group (4) contains galaxies with [NeV]/[NeII] $\geq$ 0.75 indicating an AGN contribution to the MIR greater than 50\%. Because we have usable SL spectra for only two out of the three galaxies in group (4), instead of showing the intrinsic weighted dispersion on the derived average spectra, we show the actual spectra of the two sources, normalized by the flux at 24 $\mu$m. } \label{Avg_spectra} \end{center} \end{figure} \begin{figure}[htb] \begin{center} \includegraphics[width=\linewidth]{petric_fig5.eps}%Zoom_SL.eps} \caption{Average spectra of galaxies without detectable [NeV] and with [OIV]/[NeII] flux ratios $\leq$ 0.35 (red), with detectable [NeV] emission (green), with [NeV]/[NeII] flux ratios $\geq$ 0.14 or [OIV]/[NeII] flux ratios $\geq$ 0.5 suggesting an AGN contribution to the MIR greater than $\sim$ 33\% (blue), and the average spectra of two sources in which the AGN dominates the MIR emission (magenta). } \label{Zoom-ave} \end{center} \end{figure} \begin{figure}[h] %\begin{center}$ $\begin{array}{ccccc} \includegraphics[width=0.49\linewidth]{petric_fig6a.eps}%Nov01_ne5ne2_IR.eps} %l2938 \includegraphics[width=0.49\linewidth]{petric_fig6b.eps}&%Nov01_o4ne2_IR.eps}& %l2116 \end{array}$ \caption{ Strength of the [NeV] feature (left) and [OIV] (right) emission versus the IR Luminosity as estimated from MIPS fluxes (solid symbols) or from IRAS measurements (empty symbols). } \label{HighResVsIR} \end{figure} % \begin{figure}[htb] % \begin{center} % \includegraphics[width=\linewidth]{petric_fig1b.eps} % Nov01_ne5ew_vs_ir_lum.eps % \caption{ Strength of the [NeV] feature versus the 24$\mu$m Luminosity as estimated from MIPS 24 $\mu$m measurements} % \label{eqwNe5IR} % \end{center} % \end{figure} % % \begin{figure}[htb] % \begin{center} % \includegraphics[width=\linewidth]{petric_fig2b.eps} % Nov01_o4ew_vs_ir_lum.eps % \caption{Strength of the [OIV] feature versus the 24$\mu$m Luminosity as estimated from MIPS 24 $\mu$m measurements} % \label{eqwO4IR} % \end{center} % \end{figure} % % % \begin{figure}[htb] % \begin{center} % \includegraphics[width=\linewidth]{petric_fig1c.eps} %Nov01_ne5_vs_ir_lum.eps % \caption{The ratio of the [NeV] to the [NeII] flux versus the 24$\mu$m Luminosity as estimated from MIPS 24 $\mu$m measurements} % \label{Ne5IR} % \end{center} % \end{figure} % % \begin{figure}[htb] % \begin{center} % \includegraphics[width=\linewidth]{petric_fig2c.eps} % Nov01_o4_vs_ir_lum.eps % \caption{The ratio of the [OIV] to the [NeII] flux versus the 24$\mu$m Luminosity as estimated from MIPS 24 $\mu$m measurements} % \label{O4IR} % \end{center} % \end{figure} \begin{figure}[htb] %\begin{center} $\begin{array}{cc} \includegraphics[width=0.49\linewidth]{petric_fig7a.eps}& %Nov01_pah_L24_ver1a.eps} & \includegraphics[width=0.49\linewidth]{petric_fig7b.eps}\\%Nov01_pah_L24_ver1b.eps}\\ \end{array}$ \caption{6.2 $\mu$m PAH EQW versus Log(24 $\mu$m luminosity) from MIPS 24$\mu$m estimates of the luminosity in an aperture of 1\arcmin (left) and from the nuclear LL spectra with a slit width of 10.7\arcsec (right). The 22 ULIRGs from the GOALS sample are shown in magenta filled circles, while the LIRGs are shown as black crosses. The red solid line marks where the 6.2 $\mu$m PAH EQW equals 0.53 $\mu$m. This is the average EQW for starbursts as determined by \citet{brandl06}. The dotted red lines mark the $1\sigma $ scatter in that value. While no obvious trend with luminosity can be distinguished the median 6.2 $\mu$m PAH EQW value for LIRGs is higher than that for ULIRGs. The black solid lines marks the relation between the 6.2 $\mu$m PAH EQW versus Log(24 $\mu$m luminosity) found by \citep{desai07}.} \label{Lum24-plot} %\end{center} \end{figure} \begin{figure}[htb] \begin{center} \includegraphics[width=\linewidth]{petric_fig8.eps}%Nov01_pah_f25f60.eps} \caption{FIR colors versus 6.2 $\mu$m PAH EQW for sources were MIPS 24 $\mu$m and MIPS 70 $\mu$m nuclear fluxes were extracted. The ULIRGs from the GOALS sample are shown in magenta filled circles, while the LIRGs are shown in black crosses. No significant trend with IR colors can be distinguished, however the median MIPS 24 $\mu$m and MIPS 70 $\mu$m flux ratios for AGN dominated sources are higher than those of star-formation dominated sources.} \label{FIRcolor-plot1} \end{center} \end{figure} \begin{figure}[htb] \begin{center} \includegraphics[width=\linewidth]{petric_fig9.eps}%Nov01_pah_f5f24.eps} \caption{IR colors versus 6.2 $\mu$m PAH EQW. The f5 $\mu$m and f24 $\mu$m nuclear fluxes were extracted from the low resolution spectra. The ULIRGs from the GOALS sample are shown in magenta filled circles, while the LIRGs are shown in black crosses. No significant trend with IR colors can be distinguished, however the median f5 $\mu$m to f24 $\mu$m flux ratios for AGN dominated sources are higher than those of star-formation dominated sources.} \label{FIRcolor-plot2} \end{center} \end{figure} \begin{figure}[h] %\begin{center}$ $\begin{array}{ccccc} \includegraphics[width=0.20\linewidth]{petric_fig10a.eps}&%ngc1365.jpg} & \includegraphics[width=0.21\linewidth]{petric_fig10b.eps}&%vv340.jpg}& \includegraphics[width=0.21\linewidth]{petric_fig10c.eps}&%mcg-02-01-051.jpg} & \includegraphics[width=0.21\linewidth]{petric_fig10d.eps}&%IIzw96.jpg} & \includegraphics[width=0.21\linewidth]{petric_fig10e.eps}\\%arp220.jpg} \\ \end{array}$ \caption{IRAC 3.6 $\mu$m images of NGC1365, VV340 (Armus et al. 2009), MCG-02-01-051, IIZw96 (Inami et al. 2010) and Arp 220 illustrating the merger classification stages 0 through 4 used in this paper. Stages are determined as follows(0) no obvious sign of a disturbance either in the IRAC or HST morphologies, or published evidence that the gas is not in dynamical equilibrium (i.e. undisturbed circular orbits), (1) early stage, where the galaxies are within one arc minute of each other, but little or no morphological disturbance can be observed; (2) the galaxies exhibit bridges and tidal tails but they do not have a common envelope and each optical disk is relatively intact; (3) the optical disks are completely destroyed but 2 nuclei can be distinguished; (4) the two interacting nuclei are merged but structure in the disk indicates the source has gone through a merger. } \label{mergSeq} \end{figure} \begin{figure}[htb] \begin{center} \includegraphics[width=\linewidth]{petric_fig11.eps}%Nov01_PAH_merg_nmerg.eps} \caption{Distributions of AGN dominated sources (i.e. 6.2 $\mu$m PAH EQW $\leq 0.27 \mu$m) in red and starburst objects (i.e. 6.2 $\mu$m PAH EQW $\geq 0.53 ~\mu$m and without detectable [NeV] emission) in blue as a function of merger stage (see section 4.4). Atop each bin the fraction of AGN dominated sources is written in red together with the same fraction obtained by excluding all ULIRGs (in black in parenthesis). The data is consistent with no trend between the number of AGN dominated sources and the merger stage. However the fraction of AGN dominated sources is significantly higher (40\%) for sources in the last stage of merging. If ULIRGs are excluded from the analysis, the largest fractions (26\% and 24\%) of AGN dominated LIRGs are found in the first and last merger stage respectively.} \label{merger-plot} \end{center} \end{figure} \begin{figure}[htb] \begin{center} \includegraphics[width=\linewidth]{petric_fig12.eps}%irac_plot.eps} \caption{IRAC color ([5.8 $\mu$m ] - [8 $\mu$m]) color ([3.6 $\mu$m ] - [4.5 $\mu$m]) plot of 224 LIRG nuclei. The IRAC total fluxes shown here are from apertures matched to the MIPS 24 $\mu$m apertures of typically 1', but range between 0.5' to 1.5', (Mazzarella et al. in prep.). The solid lines show the color cuts which Stern et al. (2005) use to separate active galaxies from Galactic stars and normal galaxies. The solid black, red and blue dots represent points with 6.2 $\mu$m PAH EQW greater than 0.27 $\mu$m, between 0.14 and 0.27 $\mu$m and smaller than 0.14 $\mu$m respectively. The empty circles represent nuclei without MIR spectra.} \label{irac_plot} \end{center} \end{figure} {\it{Acknowledgments:}} We thank the anonymous referee for his comments which have significantly improved our paper. VC acknowledges partial support from the EU grants ToK 39965 and FP7-REGPOT 206469. AP thanks N. Flagey for multiple readings of the document comments which helped improve and clarity of this text. AP also thanks V. Desai for help with the data analysis and C. Bridge for discussions and help with the merger classification of the LIRGs in this paper. This work is based primarily on observations made with the Spitzer Space Telescope, which is operated by the Jet Propulsion Laboratory, California Institute of Technology under NASA contract 1407. We have made use of the NASA/IPAC Extra- galactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technol- ogy, under contract with NASA. Support for this re- search was provided by NASA through an award issued by JPL/Caltech.

1012.1891

1012

1012.6027_arXiv.txt

We use the long-wavelength formalism to compute the bispectral non-Gaussianity produced in two-field inflation. We find an exact result that is used as the basis of numerical studies, and an explicit analytical slow-roll expression for several classes of potentials that gives insight into the origin and importance of the various contributions to $\f$. We also discuss the momentum dependence of $\f$. Based on these results we find a simple model that produces a relatively large non-Gaussianity. We show that the long-wavelength formalism is a viable alternative to the standard $\gd N$ formalism, and can be preferable to it in certain situations.

During inflation the energy density of the Universe is assumed to be dominated by the potential energy of one or more scalar fields in order to have a sufficiently rapid expansion to solve the homogeneity problems that plagued pre-inflationary cosmology (horizon, flatness, etc.). Very importantly, inflation also provides the initial adiabatic density perturbation that generated the large-scale structure observed today. Observations of the fluctuations in the cosmic microwave background radiation (CMB), in particular those made by the WMAP satellite, have verified the basic predictions of inflation. The most important observational parameters so far from the point of view of inflation have been the amplitude and the slope (spectral index) of the primordial power spectrum, as well as some limits on the amount of tensor perturbations and the running of the spectral index. Unfortunately this small number of observational parameters means that a very large number of quite different inflation models are all still consistent with the data. To further narrow down the number of viable inflation models additional observables are required. A very promising candidate is the non-Gaussianity of the primordial power spectrum. In its most simple form this is encoded as a non-zero three-point correlator of the CMB temperature fluctuations, or equivalently a non-zero bispectrum, which is the Fourier (or spherical harmonic on the sphere) transform of the three-point correlator. The quantity defined as the bispectrum divided by the power spectrum squared is called $\f$. The current limits on this parameter $\f$ (assuming momentum dependence of the local type, relevant for this paper) after seven years of WMAP data are $-10<\f<74$ at $95\%$ confidence level \cite{Komatsu:2010fb}. The newly launched Planck satellite \cite{:2006uk} is expected to significantly improve these constraints, down to $1\sigma$ error bars of about $3$--$5$ (depending on the use of polarization data) \cite{Komatsu:2001rj,Babich:2004yc}. While standard single-field slow-roll inflation predicts an unobservably small value of $\f$ \cite{Maldacena:2002vr}, many other models predict much larger values that could be detected or ruled out by Planck. Both supersymmetric particle theory and string theory suggest the existence of multiple scalar fields that can influence the early Universe. If more than a single scalar field plays a role during inflation, isocurvature fluctuations will be produced in addition to the adiabatic one. While these isocurvature fluctuations might have directly observable consequences in the CMB \cite{Komatsu:2010fb}, in this paper we are more interested in the effect of the isocurvature fluctuations on the adiabatic one during inflation. This effect can be important even if the isocurvature fluctuations disappear after inflation. The important point here is that while in single-field inflation the adiabatic perturbation is constant on super-horizon scales, this is no longer true in multiple-field inflation. In fact, the isocurvature perturbation acts as a source for the adiabatic perturbation on super-horizon scales and this source is multiplied by the $\eta^\perp$ parameter \cite{Rigopoulos:2005us}. This $\eta^\perp$ (defined properly in the next section) is proportional to the component of the field acceleration perpendicular to the field velocity. In other words, $\eta^\perp$ is non-zero if the field trajectory makes a turn in field space. Only during such a turn will the isocurvature mode influence the adiabatic one on super-horizon scales (see also \cite{Bernardeau:2002jy}).\footnote{This last statement is strictly only true on a flat field manifold (trivial field metric). On a curved manifold $\eta^\perp$ can be non-zero even for a straight field trajectory because of the connection terms in the covariant derivatives, see \cite{GrootNibbelink:2001qt}. In this paper a trivial field metric will be assumed.} There are two main ways to produce non-Gaussianity during inflation: during and after horizon crossing of a perturbation mode. Horizon crossing is defined as the moment when the physical wavelength $(k/a)\inv$ of the fluctuation becomes equal to the Hubble (or horizon) length $H\inv$, i.e.\ when $k=aH$ (with $k$ the wave number of the mode, $a$ the scale factor of the Universe and $H=\dot{a}/a$ the Hubble parameter). The first type of non-Gaussianity is produced in all inflation models, but it is unobservably small (i.e.\ slow-roll suppressed), unless the model contains non-standard kinetic terms (higher derivatives), like DBI inflation \cite{Alishahiha:2004eh,Langlois:2008qf} for example. In this paper we will not consider those models, but instead focus on the super-horizon type of non-Gaussianity. As is clear from the previous paragraph, super-horizon non-Gaussianity can only be produced in multiple-field inflation models where the field trajectory makes a turn in field space. To compute this type of non-Gaussianity we will make use of the long-wavelength formalism developed by Rigopoulos, Shellard and Van Tent \cite{Rigopoulos:2005xx, Rigopoulos:2005ae, Rigopoulos:2005us}, hereafter refered to as RSvT. The main purpose of this paper is twofold. In the first place we want to further work out, simplify and study the general analytic expression for $\f$ of RSvT in the case of two fields only. In the second place we want to clarify a few remaining formal issues with the formalism, and compare with an alternative formalism for computing $\f$, called the $\delta N$ formalism \cite{Starobinsky:1986fxa,Sasaki:1995aw,Sasaki:1998ug,Lyth:2004gb,Lyth:2005fi}. As will be shown, the expression for $\f$ simplifies significantly, although a final integral remains that cannot be done analytically. As a starting point for {\em numerical} work this expression is very useful though, and it gives a fully exact numerical result since no slow-roll approximation is used for the super-horizon evolution. However, to be able to derive explicit {\em analytic} results, we have to make use of the slow-roll approximation. Using this approximation we derive explicit analytic results for a number of classes of inflationary potentials. Some of these have been treated before in the literature using the $\gd N$ formalism, but others are new. In particular we show that models with a potential of the form $W(\gf,\gs) = \ga \gf^p + \gb \gs^q$ with $p=q$ will never give a large $\f$ that persists until the end of inflation (unless inflation somehow ends right during the turn of the field trajectory, but in that case a very careful treatment of the transition at the end of inflation will be required). We also present a simple model that does produce a ``large'' $\f$ of the order of a few. The reason we choose this model is that it can be treated not only numerically, but also analytically. Our analytic slow-roll results agree exactly with those derived using the $\delta N$ formalism, where available. For models where slow roll breaks down long after horizon crossing and which have to be treated numerically we also find excellent agreement. Apart from providing an alternative way of computing the bispectrum, which is always useful, the long-wavelength formalism provides a number of advantages compared to the $\gd N$ formalism. In the first place, no slow-roll approximation has to be assumed in our formalism, not even at horizon crossing. Moreover, as shown in section~\ref{Secgenmom}, it allows us to compute the momentum dependence of $\f$ due to the fact that different modes cross the horizon at different times, unlike the $\gd N$ formalism, where all modes are assumed to cross the horizon at the same time. Finally, and very importantly, the long-wavelength formalism allows for a simple physical interpretation of the different terms, showing the contributions from adiabatic and isocurvature modes and making clear why some of them can become big and others cannot. While we do not pursue this in the present paper, the formalism also provides the solution for the second-order isocurvature perturbation and hence the isocurvature bispectrum could be computed as easily as the adiabatic one. Many people have worked on non-Gaussianity, both predictions from inflation and estimators for CMB observations. The reduced bispectrum has been given in \cite{Creminelli:2005hu} for the equilateral type, in \cite{Komatsu:2001rj,Babich:2004yc} for the local type and in \cite{Senatore:2009gt} for the orthogonal type of the primordial bispectrum. The different shapes were studied in detail in \cite{Babich:2004gb,Fergusson:2008ra}. Bispectrum estimators were developed in \cite{Komatsu:2003iq,Creminelli:2005hu,Yadav:2007rk,Fergusson:2009nv, Bucher:2009nm}. All kinds of inflationary models have been studied as well. For instance, we have learned that single-field inflation models cannot produce large non-Gaussianity \cite{Maldacena:2002vr}, unless some non-trivial potential is used \cite{Chen:2006xjb} or higher derivative contributions are introduced as for the Dirac-Born-Infeld action \cite{Alishahiha:2004eh,Silverstein:2003hf,Mizuno:2009cv,Mizuno:2010ag} or K-inflation \cite{Chen:2006nt,Chen:2009bc}. The study of the effective theory of inflation has also turned out to be very fruitful \cite{Cheung:2007st}. These models were extended to incorporate multiple fields in \cite{Langlois:2008qf,Arroja:2008yy,Mizuno:2009cv,Cai:2009hw,Senatore:2010wk}. Large non-Gaussianity can also be produced at the end of inflation \cite{Lyth:2005qk,Bernardeau:2004zz,Barnaby:2006km,Enqvist:2004ey, Enqvist:2005qu,Jokinen:2005by} or after inflation, in models with varying inflaton decay rate \cite{Zaldarriaga:2003my} and in curvaton models \cite{Bartolo:2003jx,Enqvist:2005pg,Ichikawa:2008iq,Malik:2006pm,Sasaki:2006kq}. Large scale evolution of perturbations during inflation up to second order became possible through their consistent gauge invariant definition \cite{Malik:2003mv,Rigopoulos:2003ak,Langlois:2005qp}. Within the $\delta N$ formalism several authors have investigated the bispectra of specific multiple field inflation models \cite{Seery:2005gb,Kim:2006te,Battefeld:2006sz,Battefeld:2007en, Langlois:2008vk}. Two field models, being easier to deal with, have gained popularity though. Vernizzi and Wands studied the double field sum potential \cite{Vernizzi:2006ve}, while the double product potential was studied in \cite{Choi:2007su}. Conditions for large non-Gaussianity were found in \cite{Byrnes:2008wi,Byrnes:2009qy}. The paper is organized as follows. In section \ref{basic} we summarize the long-wavelength formalism of RSvT and define the various quantities used in the paper. We also give a brief overview of the $\gd N$ formalism for comparison purposes. In section \ref{general} we work out the general expression for $\f$ in the case of two fields. We also compute the momentum dependence that arises in the case that not all scales cross the horizon at the same time. These expressions, derived without using any super-horizon slow-roll approximation, are one of the main results of the paper and the starting point for our numerical analyses. In order to find completely explicit analytic expressions for $\f$, however, we do need to assume slow roll, as well as some conditions on the potential. This is treated in section \ref{secSlowRoll}, where we also compare our analytic results with those obtained using the $\gd N$ formalism, for as far as the latter exist. In section \ref{numerical} we use the two-field quadratic potential in order to compare our exact numerical results with those of the $\delta N$-formalism. We also present a simple potential that can produce an $\f$ of the order of a few, which falls into the category of potentials that can be treated analytically, thus allowing us to test our results. We conclude in section~\ref{secConclusion}. Finally, in the appendices we give supplementary information on the basis in field space that we use, comment on some gauge and formal issues, and provide several intermediate steps of our calculations. Note that \ref{appBasis} introduces in particular a small improvement of the basis defined in \cite{GrootNibbelink:2001qt} that makes it more convenient for numerical calculations during periods when the fields oscillate.

\label{secConclusion} The study of the non-Gaussianity produced by inflation models has become a hot topic of research, since the recent observations of WMAP and in particular the imminent ones of Planck will allow us to constrain and discriminate inflation models based on their non-Gaussian predictions. In this paper we investigated the super-horizon bispectral non-Gaussianity produced by two-field inflation models. To this end we further worked out the long-wavelength formalism developed by Rigopoulos, Shellard, and Van Tent (RSvT) \cite{Rigopoulos:2004ba,Rigopoulos:2005xx,Rigopoulos:2005ae, Rigopoulos:2005us}. We derived an exact result for the bispectrum parameter $\f$ produced on super-horizon scales for any two-field inflation model with canonical kinetic terms, equation~(\ref{fNLresult}). The result is expressed in terms of the linear perturbation solutions and slow-roll parameters. However, no slow-roll approximation has been assumed on super-horizon scales, these parameters should be viewed as short-hand notation and can be large. In particular this means that the result is valid for models where the field trajectory makes a sharp turn in field space so that slow roll is temporarily broken. On the other hand, we did assume slow roll to be valid at horizon crossing, but that was just for simplicity and because observations seem to indicate that this is a good approximation, it is not a fundamental assumption of the formalism. The result can be split into the sum of three parts, multiplied by an overall factor (except for a small slow-roll suppressed term that is the single-field contribution produced at horizon crossing). This overall factor is proportional to the contribution of the isocurvature mode to the adiabatic mode, which is only non-zero for a truly multiple-field model where the field trajectory makes a turn in field space, as parametrized by a non-zero value of the slow-roll parameter $\hpe$. (Effectively) single-field models do not produce any non-Gaussianity on super-horizon scales, since the adiabatic perturbation is conserved in that case. The three parts in the sum are: 1) a part that only involves slow-roll parameters evaluated at horizon-crossing and hence is always small; 2) a part proportional to the pure isocurvature mode; and 3) an integral involving terms proportional to the pure isocurvature mode. Since the adiabatic mode is not necessarily constant in the presence of isocurvature modes, we only consider models where the isocurvature mode has disappeared by the end of inflation, so that we can directly extrapolate our result at the end of inflation to recombination and observations of the CMB. However, this automatically means that the part 2), although varying wildly during the turn of the field trajectory, cannot give any persistent non-Gaussianity that can be observed in the CMB. This means that any large non-Gaussianity on super-horizon scales in models satisfying this condition will have to come from the integrated effect in part 3). The exact equation~(\ref{fNLresult}) is the basis of our numerical studies. However, to gain further insight we tried to work out the integral analytically. For this it turns out that the slow-roll approximation is necessary. Even then the integral can only be done explicitly for certain specific classes of potentials, among which are product potentials, $W(\gf,\gs) = U(\gf)V(\gs)$, and generalized sum potentials, $W(\gf,\gs) = (U(\gf)+V(\gs))^\gn$. We found that, with our assumptions on the disappearance of the isocurvature mode, no product potential can give large non-Gaussianity, nor can any simple sum potential with equal powers, $W(\gf,\gs) = \ga \gf^p + \gb \gs^p$. However, we found conditions under which the (generalized) sum potential can give large non-Gaussianity (here defined as $\f$ larger than unity), and we have described an explicit, simple model that does. It consists of a heavy field rolling down a quadratic potential while a light field sits near the local maximum of a double well potential. When the heavy field reaches zero and starts oscillating, the light field takes over and rolls down, so that there is a turn of the field trajectory in field space. We studied this model numerically, using the exact results, to confirm our analytical predictions. In deriving equation~(\ref{fNLresult}) we assumed that all three scales cross the horizon at the same moment. However, this is not a necessary assumption, and we also generalized the result to an arbitrary momentum configuration. We find that going to the squeezed limit, where one of the momenta is much smaller than the other two, even when remaining within the resolution of the Planck satellite ($k' \sim 1000 k$), the result for $\f$ can be reduced by about $10\%$. We stress that we are discussing $\f$ here, so this effect is unrelated to the well-known result that the local bispectrum peaks on squeezed momentum configurations, which is due to the momentum behaviour of the power spectrum, which has been divided out in $\f$. However, exactly because of this latter effect, the squeezed limit is very relevant for the computation of $\f$. We managed to work out and include the second-order source term at horizon crossing in the long-wavelength formalism of RSvT, a contribution that had been missing so far. While this term is always small for the models we consider, it means we could compare our results directly to the so-called $\f^{(4)}$ as defined in the $\gd N$ formalism \cite{Choi:2007su,Byrnes:2008wi,Vernizzi:2006ve}. Some of the potentials we studied had already been worked out using that formalism and where available we compared our analytical results and found perfect agreement. We also compared our exact numerical results with those obtained using a numerical $\gd N$ treatment for models where slow roll is broken and the analytic results cannot be trusted, and again we found excellent agreement. We showed that the long-wavelength formalism of RSvT represents a viable alternative to the $\gd N$ formalism to compute the super-horizon non-Gaussianity produced during inflation, allowing us to obtain and verify results in a different way. Moreover, the long-wavelength formalism has a number of advantages that make it preferable in certain situations. In the first place, unlike for the $\gd N$ formalism, the assumption of slow roll at horizon crossing is not fundamental, one can use the exact numerical linear solutions instead. Secondly, the long-wavelength formalism does not require the assumption that all scales cross the horizon at the same time, unlike the $\gd N$ formalism. Since the squeezed limit is most relevant for super-horizon non-Gaussianity, the ability of our formalism to take this momentum-dependent correction of $\f$ into account can be important, depending on the model. Finally, our formalism allows for a simple physical interpretation of the different parts in terms of adiabatic and isocurvature modes, providing insight into the behaviour of the different transient and persistent contributions to $\f$. From our studies it has become clear that the condition on the disappearance of the isocurvature mode by the end of inflation is a very strong constraint. It significantly reduces the possibilities for a large, observable value of $\f$ produced during inflation. In future work we would like to relax this condition. However, since that means the adiabatic mode is no longer necessarily constant after inflation, this will require a much better description and understanding of the evolution of the perturbations during the transition at the end of inflation and the subsequent period of (p)reheating. In conclusion, while a lot of progress has been made over the past few years regarding the non-Gaussianity produced in multiple-field inflation, more work still remains to be done. \ack BvT would like to thank Gerasimos Rigopoulos and Paul Shellard for many useful discussions, especially in the initial stages of this work. The authors would also like to thank Filippo Vernizzi for useful discussions. \appendix \label{app}

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{ We propose a new way to implement an inflationary prior to a cosmological dataset that incorporates the inflationary observables at arbitrary order. This approach employs an exponential form for the Hubble parameter $H(\phi)$ without taking the slow-roll approximation. At lowest non-trivial order, this $H(\phi)$ has the unique property that it is the solution to the brachistochrone problem for inflation. } \begin{document}

In inflationary universe, we can use the evolution of the inflaton as a clock. For a single scalar field with canonical kinetic term, the inflationary Hubble parameter can be expressed as a function of $\phi$ only: \begin{equation} H(\phi, \dot\phi) \rightarrow H(\phi) \end{equation} Since we need an almost flat inflaton potential for enough e-folds of inflation, we may choose to start with a truncated expression \cite{inflation_flow}, for some $M$, \begin{equation} \label{Hflat} H(\phi) = H_0 (1 + b_1 \f + b_2 \f^2/2 + \dots + b_M \f^M/M!) \end{equation} where $\phi$ is in Planck units and $b_i$'s are dimensionless. The coefficients $b_i$ are expected to be small, so the truncation to $M+1$ terms, with a relatively small $M$, should yield a very good approximation, especially for the region where CMB and other cosmological data are available to constrain them. Most efforts to implement an inflationary prior to a cosmological data set along this line use $H(\phi)$ of this form \cite{Hansen:2001eu, Kinney:2003uw, Peiris:2006sj}. An alternative approach uses the inflationary flow formalism to reconstruct models without assuming slow roll \cite{Powell:2007gu}. One might worry that different parameterizations would lead to different constraints on the inflationary parameters. It was shown in \cite{Hamann:2008pb} that third-order Taylor expansions of $H(\phi)$ and $H^2(\phi)$ with a prior of sufficient e-folds yield the same constraints on the observable window of inflation. The conditions for a reasonable trial $H(\phi)$ are: first, it is close to what data indicate, so the Taylor expansion needs only a few terms; second, the expression is easy to manipulate when calculating observables. Third, it will be nice if the leading term has a simple physical interpretation. In the above case, the leading term is simply a flat potential yielding an unlimited number of e-folds. $H(\phi)$ of Eq.(\ref{Hflat}) is just a perturbed version of a flat inflaton potential. We like to ask if there is another expression for $H(\phi)$ which has these three properties. Furthermore, it will be really useful if one does not have to take the slow-roll approximation. Here we claim that there is such an expression, which in addition has a conceptual underpinning behind it. Consider \begin{equation} \label{form1} H^{(M)}(\phi) = H_0 \exp(a_1\f + a_2\f^2/2 + a_3 \f^3/3! + \dots + a_M \f^M/M!) ~, \end{equation} here $H^{(0)}(\phi)=H_0$ corresponds to a flat potential. What is particularly interesting about this exponential form is that, when it is truncated to the linear term in the exponent, i.e., \begin{equation} \label{brach} H^{(1)}(\phi) = H_0\exp(a_1 \phi) = H_0 \exp(\sqrt{\epsilon /2}\, \phi) \end{equation} it is simply the solution to the brachistochrone problem for inflation without taking the slow-roll approximation. That is, away from the slow-roll approximation, $H^{(1)}(\phi)$ yields the minimum number of e-folds (fastest path) for a fixed drop in $H$ over a fixed field range. So, any deviation from this path will yield more e-folds. In this sense, it is the opposite of the flat potential. The exponential form of $H(\phi)$ was also mentioned in an early paper \cite{Liddle:1993ch} by Andrew Liddle, as a way to construct arbitrary inflaton potential. Here, we motivate the form Eq.(\ref{form1}) from entirely new perspective, i.e. the brachistochrone problem for inflation and how to deviate away from the brachistochrone solution. In this sense, we have given the exponential of $H(\phi)$ an interesting physical interpretation. Furthermore, it is quite amazing that, for small $\epsilon$, $H^{(1)}(\phi)$ alone (i.e., without higher terms in Eq.(\ref{form1})) can yield an inflationary scenario within experimental bounds, implying a low (small $M$, say $M \gtrsim 2$) $H^{(M)}(\phi)$ will do very well already. Models that are close to the brachistochrone solution typically require a spectral index that is within a few percent of unity. Lastly, the relations between the slow-roll parameters and the coefficients $a_i$ in $H^{(M)}(\phi)$ of Eq.(\ref{form1}) are very simple.

In this paper, we have proposed a new functional form of $H(\phi)$ as an effective way to implement inflationary priors in cosmological data analysis. This new form of $H(\phi)$ entails writing $H$ in the exponential form. At the lowest order, with the exponent linear in $\phi$, it is the solution to the brachistochrone problem for inflation, which corresponds to the minimal number of e-folds for a fixed drop in $H$ over a fixed field range. Higher-order terms can be included to deviate the inflaton path away from the brachistochrone solution. In addition to the theoretical underpinning, the exponential form of $H(\phi)$ also provides an efficient way to compute power spectrum observables. If one Taylor expands the function $\ln (H(\phi))$, the expansion coefficients are natural parameters to express the observables, such as $n_s$, $\ud n_s /\ud \ln k$, $r$ and $n_t$. Higher-order derivatives of the power spectrum entails higher-order expansion coefficients of $\ln (H(\phi))$, and the computation is more straightforward than the usual $\epsilon$ and $\eta$ parametrization. We have also performed MCMC analysis to illustrate how observation data can constrain the expansion coefficients of $\ln (H(\phi))$. The WMAP7 data constrains the two leading coefficients very well, but is not quite sensitive to the coefficients of the cubic and quartic terms. We expect including more data on different scales will improve the constraints. With the upcoming precision measurements from Cosmic Microwave Background and Large Scale Structure, we hope our proposal will offer an efficient way to reconstruct the inflaton potential from data. In the actual implementation of an inflationary prior to a cosmological dataset, one does not need to use the same $H(\phi)$ for the whole range of sixty e-folds of inflation. The actual inflaton potential can be more complicated; in fact, it can be multi-field. In these more general situations, one can use a set of piecewise exponential segments of the form (\ref{form1}) instead. In this case, the constraint of sixty or more e-folds imposed in the above analysis may be relaxed. \vspace{0.6cm} \noindent {\bf Acknowledgments} \vspace{0.3cm} DW thanks N. Agarwal for technical assistance. DW and HT are supported by the National Science Foundation under grant PHY-0355005. JX was supported in part by a DOE grant DE-FG-02-95ER40896, a Cottrell Scholar Award from Research Corporation, and a Vilas Associate Award. \vspace{0.5cm}

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